GNU Astronomy Utilities

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GNU Astronomy Utilities

This book documents version 0.7 of the GNU Astronomy Utilities (Gnuastro). Gnuastro provides various programs and libraries for astronomical data manipulation and analysis.

Copyright © 2015-2018 Free Software Foundation, Inc.

Permission is granted to copy, distribute and/or modify this document under the terms of the GNU Free Documentation License, Version 1.3 or any later version published by the Free Software Foundation; with no Invariant Sections, no Front-Cover Texts, and no Back-Cover Texts. A copy of the license is included in the section entitled “GNU Free Documentation License”.

To navigate easily in this web page, you can use the Next, Previous, Up and Contents links in the top and bottom of each page. Next and Previous will take you to the next or previous topic in the same level, for example from chapter 1 to chapter 2 or vice versa. To go to the sections or subsections, you have to click on the menu entries that are there when ever a sub-component to a title is present.

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1 Introduction

GNU Astronomy Utilities (Gnuastro) is an official GNU package consisting of separate programs and libraries for the manipulation and analysis of astronomical data. All the programs share the same basic command-line user interface for the comfort of both the users and developers. Gnuastro is written to comply fully with the GNU coding standards so it integrates finely with the GNU/Linux operating system. This also enables astronomers to expect a fully familiar experience in the source code, building, installing and command-line user interaction that they have seen in all the other GNU software that they use. The official and always up to date version of this book (or manual) is freely available under GNU Free Doc. License in various formats (PDF, HTML, plain text, info, and as its Texinfo source) at

For users who are new to the GNU/Linux environment, unless otherwise specified most of the topics in Installation and Common program behavior are common to all GNU software, for example installation, managing command-line options or getting help (also see New to GNU/Linux?). So if you are new to this empowering environment, we encourage you to go through these chapters carefully. They can be a starting point from which you can continue to learn more from each program’s own manual and fully benefit from and enjoy this wonderful environment. Gnuastro also comes with a large set of libraries, so you can write your own programs using Gnuastro’s building blocks, see Review of library fundamentals for an introduction.

In Gnuastro, no change to any program or library will be committed to its history, before it has been fully documented here first. As discussed in Science and its tools this is a founding principle of the Gnuastro.

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1.1 Quick start

The latest official release tarball is always available as gnuastro-latest.tar.gz. For better compression (faster download), and robust archival features, an Lzip compressed tarball is also available at gnuastro-latest.tar.lz, see Release tarball for more details on the tarball release1.

Let’s assume the downloaded tarball is in the TOPGNUASTRO directory. The first two commands below can be used to decompress the source. If you download tar.lz and your Tar implementation doesn’t recognize Lzip (the second command fails), run the third and fourth lines2. Note that lines starting with ## don’t need to be typed.

## Go into the download directory.

## Also works on `tar.gz'. GNU Tar recognizes both formats.
$ tar xf gnuastro-latest.tar.lz

## Only when previous command fails.
$ lzip -d gnuastro-latest.tar.lz
$ tar xf gnuastro-latest.tar

Gnuastro has three mandatory dependencies and some optional dependencies for extra functionality, see Dependencies for the full list. In Dependencies from package managers we have prepared the command to easily install Gnuastro’s dependencies using the package manager of some operating systems. When the mandatory dependencies are ready, you can configure, compile, check and install Gnuastro on your system with the following commands.

$ cd gnuastro-X.X                  # Replace X.X with version number.
$ ./configure
$ make -j8                         # Replace 8 with no. CPU threads.
$ make check
$ sudo make install

See Known issues if you confront any complications. For each program there is an ‘Invoke ProgramName’ sub-section in this book which explains how the programs should be run on the command-line (for example Invoking Table). You can read the same section on the command-line by running $ info astprogname (for example info asttable). The ‘Invoke ProgramName’ sub-section starts with a few examples of each program and goes on to explain the invocation details. See Getting help for all the options you have to get help. In Tutorials some real life examples of how these programs might be used are given.

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1.2 Science and its tools

History of science indicates that there are always inevitably unseen faults, hidden assumptions, simplifications and approximations in all our theoretical models, data acquisition and analysis techniques. It is precisely these that will ultimately allow future generations to advance the existing experimental and theoretical knowledge through their new solutions and corrections.

In the past, scientists would gather data and process them individually to achieve an analysis thus having a much more intricate knowledge of the data and analysis. The theoretical models also required little (if any) simulations to compare with the data. Today both methods are becoming increasingly more dependent on pre-written software. Scientists are dissociating themselves from the intricacies of reducing raw observational data in experimentation or from bringing the theoretical models to life in simulations. These ‘intricacies’ are precisely those unseen faults, hidden assumptions, simplifications and approximations that define scientific progress.

Unfortunately, most persons who have recourse to a computer for statistical analysis of data are not much interested either in computer programming or in statistical method, being primarily concerned with their own proper business. Hence the common use of library programs and various statistical packages. ... It’s time that was changed.

F. J. Anscombe. The American Statistician, Vol. 27, No. 1. 1973

Anscombe’s quartet demonstrates how four data sets with widely different shapes (when plotted) give nearly identical output from standard regression techniques. Anscombe uses this (now famous) quartet, which was introduced in the paper quoted above, to argue that “Good statistical analysis is not a purely routine matter, and generally calls for more than one pass through the computer”. Echoing Anscombe’s concern after 44 years, some of the highly recognized statisticians of our time (Leek, McShane, Gelman, Colquhoun, Nuijten and Goodman), wrote in Nature that:

We need to appreciate that data analysis is not purely computational and algorithmic — it is a human behaviour....Researchers who hunt hard enough will turn up a result that fits statistical criteria — but their discovery will probably be a false positive.

Five ways to fix statistics, Nature, 551, Nov 2017.

Users of statistical (scientific) methods (software) are therefore not passive (objective) agents in their result. Therefore, it is necessary to actually understand the method, not just use it as a black box. The subjective experience gained by frequently using a method/software is not sufficient to claim an understanding of how the tool/method works and how relevant it is to the data and analysis. This kind of subjective experience is prone to serious misunderstandings about the data, what the software/statistical-method really does (especially as it gets more complicated), and thus the scientific interpretation of the result. This attitude is further encouraged through non-free software3, poorly written (or non-existent) scientific software manuals, and non-reproducible papers4. This approach to scientific software and methods only helps in producing dogmas and an “obscurantist faith in the expert’s special skill, and in his personal knowledge and authority5.

Program or be programmed. Choose the former, and you gain access to the control panel of civilization. Choose the latter, and it could be the last real choice you get to make.

Douglas Rushkoff. Program or be programmed, O/R Books (2010).

It is obviously impractical for any one human being to gain the intricate knowledge explained above for every step of an analysis. On the other hand, scientific data can be large and numerous, for example images produced by telescopes in astronomy. This requires efficient algorithms. To make things worse, natural scientists have generally not been trained in the advanced software techniques, paradigms and architecture that are taught in computer science or engineering courses and thus used in most software. The GNU Astronomy Utilities are an effort to tackle this issue.

Gnuastro is not just a software, this book is as important to the idea behind Gnuastro as the source code (software). This book has tried to learn from the success of the “Numerical Recipes” book in educating those who are not software engineers and computer scientists but still heavy users of computational algorithms, like astronomers. There are two major differences.

The first difference is that Gnuastro’s code and the background information are segregated: the code is moved within the actual Gnuastro software source code and the underlying explanations are given here in this book. In the source code, every non-trivial step is heavily commented and correlated with this book, it follows the same logic of this book, and all the programs follow a similar internal data, function and file structure, see Program source. Complementing the code, this book focuses on thoroughly explaining the concepts behind those codes (history, mathematics, science, software and usage advise when necessary) along with detailed instructions on how to run the programs. At the expense of frustrating “professionals” or “experts”, this book and the comments in the code also intentionally avoid jargon and abbreviations. The source code and this book are thus intimately linked, and when considered as a single entity can be thought of as a real (an actual software accompanying the algorithms) “Numerical Recipes” for astronomy.

The second major, and arguably more important, difference is that “Numerical Recipes” does not allow you to distribute any code that you have learned from it. In other words, it does not allow you to release your software’s source code if you have used their codes, you can only publicly release binaries (a black box) to the community. Therefore, while it empowers the privileged individual who has access to it, it exacerbates social ignorance. Exactly at the opposite end of the spectrum, Gnuastro’s source code is released under the GNU general public license (GPL) and this book is released under the GNU free documentation license. You are therefore free to distribute any software you create using parts of Gnuastro’s source code or text, or figures from this book, see Your rights.

With these principles in mind, Gnuastro’s developers aim to impose the minimum requirements on you (in computer science, engineering and even the mathematics behind the tools) to understand and modify any step of Gnuastro if you feel the need to do so, see Why C programming language? and Program design philosophy.

Without prior familiarity and experience with optics, it is hard to imagine how, Galileo could have come up with the idea of modifying the Dutch military telescope optics to use in astronomy. Astronomical objects could not be seen with the Dutch military design of the telescope. In other words, it is unlikely that Galileo could have asked a random optician to make modifications (not understood by Galileo) to the Dutch design, to do something no astronomer of the time took seriously. In the paradigm of the day, what could be the purpose of enlarging geometric spheres (planets) or points (stars)? In that paradigm only the position and movement of the heavenly bodies was important, and that had already been accurately studied (recently by Tycho Brahe).

In the beginning of his “The Sidereal Messenger” (published in 1610) he cautions the readers on this issue and before describing his results/observations, Galileo instructs us on how to build a suitable instrument. Without a detailed description of how he made his tools and done his observations, no reasonable person would believe his results. Before he actually saw the moons of Jupiter, the mountains on the Moon or the crescent of Venus, Galileo was “evasive”6 to Kepler. Science is defined by its tools/methods, not its results7

The same is true today: science cannot progress with a black box, or poorly released code. Technical knowledge and experience (to experiment on its tools, or software in this context8), is critical to scientific vitality. Scientific research are only considered for peer review and publication if they have a sufficiently high standard of English style. A similar level of quality assessment is necessary regarding the codes/methods scientists use to derive their results. Therefore, when a scientist says “software is not my specialty, I am not a software engineer. So the quality of my code/processing doesn’t matter. Why should I master good coding style, or release my code, when I am hired to do Astronomy/Biology?”. This statement is akin to a French scientist saying that "English is not my language, I am not Shakespeare. So the quality of my English writing doesn’t matter. Why should I master good English style, when I am hired to do Astronomy/Biology?"

Bjarne Stroustrup (creator of the C++ language) says: “Without understanding software, you are reduced to believing in magic”. Ken Thomson (the designer or the Unix operating system) says “I abhor a system designed for the ‘user’ if that word is a coded pejorative meaning ‘stupid and unsophisticated’.” Certainly no scientist (user of a scientific software) would want to be considered a believer in magic, or stupid and unsophisticated.

This can happen when scientists get too distant from the raw data and methods, and are mainly discussing results. In other words, when they feel they have tamed Nature into their own high-level (abstract) models (creations), and are mainly concerned with scaling up, or industrializing those results. Roughly five years before special relativity, and about two decades before quantum mechanics fundamentally changed Physics, Lord Kelvin is quoted as saying:

There is nothing new to be discovered in physics now. All that remains is more and more precise measurement.

William Thomson (Lord Kelvin), 1900

A few years earlier Albert. A. Michelson made the following statement:

The more important fundamental laws and facts of physical science have all been discovered, and these are now so firmly established that the possibility of their ever being supplanted in consequence of new discoveries is exceedingly remote.... Our future discoveries must be looked for in the sixth place of decimals.

Albert. A. Michelson, dedication of Ryerson Physics Lab, U. Chicago 1894

If scientists are considered to be more than mere “puzzle” solvers9 (simply adding to the decimals of existing values or observing a feature in 10, 100, or 100000 more galaxies or stars, as Kelvin and Michelson clearly believed), they cannot just passively sit back and uncritically repeat the previous (observational or theoretical) methods/tools on new data. Today there is a wealth of raw telescope images ready (mostly for free) at the finger tips of anyone who is interested with a fast enough internet connection to download them. The only thing lacking is new ways to analyze this data and dig out the treasure that is lying hidden in them to existing methods and techniques.

New data that we insist on analyzing in terms of old ideas (that is, old models which are not questioned) cannot lead us out of the old ideas. However many data we record and analyze, we may just keep repeating the same old errors, missing the same crucially important things that the experiment was competent to find.

Jaynes, Probability theory, the logic of science. Cambridge U. Press (2003).

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1.3 Your rights

The paragraphs below, in this section, belong to the GNU Texinfo10 manual and are not written by us! The name “Texinfo” is just changed to “GNU Astronomy Utilities” or “Gnuastro” because they are released under the same licenses and it is beautifully written to inform you of your rights.

GNU Astronomy Utilities is “free software”; this means that everyone is free to use it and free to redistribute it on certain conditions. Gnuastro is not in the public domain; it is copyrighted and there are restrictions on its distribution, but these restrictions are designed to permit everything that a good cooperating citizen would want to do. What is not allowed is to try to prevent others from further sharing any version of Gnuastro that they might get from you.

Specifically, we want to make sure that you have the right to give away copies of the programs that relate to Gnuastro, that you receive the source code or else can get it if you want it, that you can change these programs or use pieces of them in new free programs, and that you know you can do these things.

To make sure that everyone has such rights, we have to forbid you to deprive anyone else of these rights. For example, if you distribute copies of the Gnuastro related programs, you must give the recipients all the rights that you have. You must make sure that they, too, receive or can get the source code. And you must tell them their rights.

Also, for our own protection, we must make certain that everyone finds out that there is no warranty for the programs that relate to Gnuastro. If these programs are modified by someone else and passed on, we want their recipients to know that what they have is not what we distributed, so that any problems introduced by others will not reflect on our reputation.

The full text of the licenses for the Gnuastro book and software can be respectively found in GNU Gen. Pub. License v311 and GNU Free Doc. License12.

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1.4 Naming convention

Gnuastro is a package of independent programs and a collection of libraries, here we are mainly concerned with the programs. Each program has an official name which consists of one or two words, describing what they do. The latter are printed with no space, for example NoiseChisel or Crop. On the command-line, you can run them with their executable names which start with an ast and might be an abbreviation of the official name, for example astnoisechisel or astcrop, see Executable names.

We will use “ProgramName” for a generic official program name and astprogname for a generic executable name. In this book, the programs are classified based on what they do and thoroughly explained. An alphabetical list of the programs that are installed on your system with this installation are given in Gnuastro programs list. That list also contains the executable names and version numbers along with a one line description.

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1.5 Version numbering

Gnuastro can have two formats of version numbers, for official and unofficial releases. Official Gnuastro releases are announced on the info-gnuastro mailing list, they have a version control tag in Gnuastro’s development history, and their version numbers are formatted like “A.B”. A is a major version number, marking a significant planned achievement (for example see GNU Astronomy Utilities 1.0), while B is a minor version number, see below for more on the distinction. Note that the numbers are not decimals, so version 2.34 is much more recent than version 2.5, which is not equal to 2.50.

Gnuastro also allows a unique version number for unofficial releases. Unofficial releases can mark any point in Gnuastro’s development history. This is done to allow astronomers to easily use any point in the version controlled history for their data-analysis and research publication. See Version controlled source for a complete introduction. This section is not just for developers and is intended to straightforward and easy to read, so please have a look if you are interested in the cutting-edge. This unofficial version number is a meaningful and easy to read string of characters, unique to that particular point of history. With this feature, users can easily stay up to date with the most recent bug fixes and additions that are committed between official releases.

The unofficial version number is formatted like: A.B.C-D. A and B are the most recent official version number. C is the number of commits that have been made after version A.B. D is the first 4 or 5 characters of the commit hash number13. Therefore, the unofficial version number ‘3.92.8-29c8’, corresponds to the 8th commit after the official version 3.92 and its commit hash begins with 29c8. The unofficial version number is sort-able (unlike the raw hash) and as shown above is descriptive of the state of the unofficial release. Of course an official release is preferred for publication (since its tarballs are easily available and it has gone through more tests, making it more stable), so if an official release is announced prior to your publication’s final review, please consider updating to the official release.

The major version number is set by a major goal which is defined by the developers and user community before hand, for example see GNU Astronomy Utilities 1.0. The incremental work done in minor releases are commonly small steps in achieving the major goal. Therefore, there is no limit on the number of minor releases and the difference between the (hypothetical) versions 2.927 and 3.0 can be a small (negligible to the user) improvement that finalizes the defined goals.

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1.5.1 GNU Astronomy Utilities 1.0

Currently (prior to Gnuastro 1.0), the aim of Gnuastro is to have a complete system for data manipulation and analysis at least similar to IRAF14. So an astronomer can take all the standard data analysis steps (starting from raw data to the final reduced product and standard post-reduction tools) with the various programs in Gnuastro.

The maintainers of each camera or detector on a telescope can provide a completely transparent shell script or Makefile to the observer for data analysis. This script can set configuration files for all the required programs to work with that particular camera. The script can then run the proper programs in the proper sequence. The user/observer can easily follow the standard shell script to understand (and modify) each step and the parameters used easily. Bash (or other modern GNU/Linux shell scripts) is powerful and made for this gluing job. This will simultaneously improve performance and transparency. Shell scripting (or Makefiles) are also basic constructs that are easy to learn and readily available as part of the Unix-like operating systems. If there is no program to do a desired step, Gnuastro’s libraries can be used to build specific programs.

The main factor is that all observatories or projects can freely contribute to Gnuastro and all simultaneously benefit from it (since it doesn’t belong to any particular one of them), much like how for-profit organizations (for example RedHat, or Intel and many others) are major contributors to free and open source software for their shared benefit. Gnuastro’s copyright has been fully awarded to GNU, so it doesn’t belong to any particular astronomer or astronomical facility or project.

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1.6 New to GNU/Linux?

Some astronomers initially install and use a GNU/Linux operating system because their necessary tools can only be installed in this environment. However, the transition is not necessarily easy. To encourage you in investing the patience and time to make this transition, and actually enjoy it, we will first start with a basic introduction to GNU/Linux operating systems. Afterwards, in Command-line interface we’ll discuss the wonderful benefits of the command-line interface, how it beautifully complements the graphic user interface, and why it is worth the (apparently steep) learning curve. Finally a complete chapter (Tutorials) is devoted to real world scenarios of using Gnuastro (on the command-line). Therefore if you don’t yet feel comfortable with the command-line we strongly recommend going through that chapter after finishing this section.

You might have already noticed that we are not using the name “Linux”, but “GNU/Linux”. Please take the time to have a look at the following essays and FAQs for a complete understanding of this very important distinction.

In short, the Linux kernel15 is built using the GNU C library (glibc) and GNU compiler collection (gcc). The Linux kernel software alone is just a means for other software to access the hardware resources, it is useless alone: to say “running Linux”, is like saying “driving your carburetor”.

To have an operating system, you need lower-level (to build the kernel), and higher-level (to use it) software packages. The majority of such software in most Unix-like operating systems are GNU software: “the whole system is basically GNU with Linux loaded”. Therefore to acknowledge GNU’s instrumental role in the creation and usage of the Linux kernel and the operating systems that use it, we should call these operating systems “GNU/Linux”.

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1.6.1 Command-line interface

One aspect of Gnuastro that might be a little troubling to new GNU/Linux users is that (at least for the time being) it only has a command-line user interface (CLI). This might be contrary to the mostly graphical user interface (GUI) experience with proprietary operating systems. Since the various actions available aren’t always on the screen, the command-line interface can be complicated, intimidating, and frustrating for a first-time user. This is understandable and also experienced by anyone who started using the computer (from childhood) in a graphical user interface (this includes most of Gnuastro’s authors). Here we hope to convince you of the unique benefits of this interface which can greatly enhance your productivity while complementing your GUI experience.

Through GNOME 316, most GNU/Linux based operating systems now have an advanced and useful GUI. Since the GUI was created long after the command-line, some wrongly consider the command line to be obsolete. Both interfaces are useful for different tasks. For example you can’t view an image, video, pdf document or web page on the command-line. On the other hand you can’t reproduce your results easily in the GUI. Therefore they should not be regarded as rivals but as complementary user interfaces, here we will outline how the CLI can be useful in scientific programs.

You can think of the GUI as a veneer over the CLI to facilitate a small subset of all the possible CLI operations. Each click you do on the GUI, can be thought of as internally running a different CLI command. So asymptotically (if a good designer can design a GUI which is able to show you all the possibilities to click on) the GUI is only as powerful as the command-line. In practice, such graphical designers are very hard to find for every program, so the GUI operations are always a subset of the internal CLI commands. For programs that are only made for the GUI, this results in not including lots of potentially useful operations. It also results in ‘interface design’ to be a crucially important part of any GUI program. Scientists don’t usually have enough resources to hire a graphical designer, also the complexity of the GUI code is far more than CLI code, which is harmful for a scientific software, see Science and its tools.

For programs that have a GUI, one action on the GUI (moving and clicking a mouse, or tapping a touchscreen) might be more efficient and easier than its CLI counterpart (typing the program name and your desired configuration). However, if you have to repeat that same action more than once, the GUI will soon become frustrating and prone to errors. Unless the designers of a particular program decided to design such a system for a particular GUI action, there is no general way to run any possible series of actions automatically on the GUI.

On the command-line, you can run any series of of actions which can come from various CLI capable programs you have decided your self in any possible permutation with one command17. This allows for much more creativity and exact reproducibility that is not possible to a GUI user. For technical and scientific operations, where the same operation (using various programs) has to be done on a large set of data files, this is crucially important. It also allows exact reproducibility which is a foundation principle for scientific results. The most common CLI (which is also known as a shell) in GNU/Linux is GNU Bash, we strongly encourage you to put aside several hours and go through this beautifully explained web page: You don’t need to read or even fully understand the whole thing, only a general knowledge of the first few chapters are enough to get you going.

Since the operations in the GUI are limited and they are visible, reading a manual is not that important in the GUI (most programs don’t even have any!). However, to give you the creative power explained above, with a CLI program, it is best if you first read the manual of any program you are using. You don’t need to memorize any details, only an understanding of the generalities is needed. Once you start working, there are more easier ways to remember a particular option or operation detail, see Getting help.

To experience the command-line in its full glory and not in the GUI terminal emulator, press the following keys together: CTRL+ALT+F418 to access the virtual console. To return back to your GUI, press the same keys above replacing F4 with F7 (or F1, or F2, depending on your GNU/Linux distribution). In the virtual console, the GUI, with all its distracting colors and information, is gone. Enabling you to focus entirely on your actual work.

For operations that use a lot of your system’s resources (processing a large number of large astronomical images for example), the virtual console is the place to run them. This is because the GUI is not competing with your research work for your system’s RAM and CPU. Since the virtual consoles are completely independent, you can even log out of your GUI environment to give even more of your hardware resources to the programs you are running and thus reduce the operating time.

Since it uses far less system resources, the CLI is also convenient for remote access to your computer. Using secure shell (SSH) you can log in securely to your system (similar to the virtual console) from anywhere even if the connection speeds are low. There are apps for smart phones and tablets which allow you to do this.

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1.7 Report a bug

According to Wikipedia “a software bug is an error, flaw, failure, or fault in a computer program or system that causes it to produce an incorrect or unexpected result, or to behave in unintended ways”. So when you see that a program is crashing, not reading your input correctly, giving the wrong results, or not writing your output correctly, you have found a bug. In such cases, it is best if you report the bug to the developers. The programs will also inform you if known impossible situations occur (which are caused by something unexpected) and will ask the users to report the bug issue.

Prior to actually filing a bug report, it is best to search previous reports. The issue might have already been found and even solved. The best place to check if your bug has already been discussed is the bugs tracker on Gnuastro project webpage at In the top search fields (under “Display Criteria”) set the “Open/Closed” drop-down menu to “Any” and choose the respective program or general category of the bug in “Category” and click the “Apply” button. The results colored green have already been solved and the status of those colored in red is shown in the table.

Recently corrected bugs are probably not yet publicly released because they are scheduled for the next Gnuastro stable release. If the bug is solved but not yet released and it is an urgent issue for you, you can get the version controlled source and compile that, see Version controlled source.

To solve the issue as readily as possible, please follow the following to guidelines in your bug report. The How to Report Bugs Effectively and How To Ask Questions The Smart Way essays also provide some good generic advice for all software (don’t contact their authors for Gnuastro’s problems). Mastering the art of giving good bug reports (like asking good questions) can greatly enhance your experience with any free and open source software. So investing the time to read through these essays will greatly reduce your frustration after you see something doesn’t work the way you feel it is supposed to for a large range of software, not just Gnuastro.

Be descriptive

Please provide as many details as possible and be very descriptive. Explain what you expected and what the output was: it might be that your expectation was wrong. Also please clearly state which sections of the Gnuastro book (this book), or other references you have studied to understand the problem. This can be useful in correcting the book (adding links to likely places where users will check). But more importantly, it will be encouraging for the developers, since you are showing how serious you are about the problem and that you have actually put some thought into it. “To be able to ask a question clearly is two-thirds of the way to getting it answered.” – John Ruskin (1819-1900).

Individual and independent bug reports

If you have found multiple bugs, please send them as separate (and independent) bugs (as much as possible). This will significantly help us in managing and resolving them sooner.

Reproducible bug reports

If we cannot exactly reproduce your bug, then it is very hard to resolve it. So please send us a Minimal working example19 along with the description. For example in running a program, please send us the full command-line text and the output with the -P option, see Operating mode options. If it is caused only for a certain input, also send us that input file. In case the input FITS is large, please use Crop to only crop the problematic section and make it as small as possible so it can easily be uploaded and downloaded and not waste the archive’s storage, see Crop.

There are generally two ways to inform us of bugs:

Once the items have been registered in the mailing list or webpage, the developers will add it to either the “Bug Tracker” or “Task Manager” trackers of the Gnuastro project webpage. These two trackers can only be edited by the Gnuastro project developers, but they can be browsed by anyone, so you can follow the progress on your bug. You are most welcome to join us in developing Gnuastro and fixing the bug you have found maybe a good starting point. Gnuastro is designed to be easy for anyone to develop (see Science and its tools) and there is a full chapter devoted to developing it: Developing.

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1.8 Suggest new feature

We would always be happy to hear of suggested new features. For every program there are already lists of features that we are planning to add. You can see the current list of plans from the Gnuastro project webpage at and following “Tasks”→“Browse” on the horizontal menu at the top of the page immediately under the title, see Gnuastro project webpage. If you want to request a feature to an existing program, click on the “Display Criteria” above the list and under “Category”, choose that particular program. Under “Category” you can also see the existing suggestions for new programs or other cases like installation, documentation or libraries. Also be sure to set the “Open/Closed” value to “Any”.

If the feature you want to suggest is not already listed in the task manager, then follow the steps that are fully described in Report a bug. Please have in mind that the developers are all busy with their own astronomical research, and implementing existing “task”s to add or resolving bugs. Gnuastro is a volunteer effort and none of the developers are paid for their hard work. So, although we will try our best, please don’t not expect that your suggested feature be immediately included (with the next release of Gnuastro).

The best person to apply the exciting new feature you have in mind is you, since you have the motivation and need. In fact Gnuastro is designed for making it as easy as possible for you to hack into it (add new features, change existing ones and so on), see Science and its tools. Please have a look at the chapter devoted to developing (Developing) and start applying your desired feature. Once you have added it, you can use it for your own work and if you feel you want others to benefit from your work, you can request for it to become part of Gnuastro. You can then join the developers and start maintaining your own part of Gnuastro. If you choose to take this path of action please contact us before hand (Report a bug) so we can avoid possible duplicate activities and get interested people in contact.

Gnuastro is a collection of low level programs: As described in Program design philosophy, a founding principle of Gnuastro is that each library or program should be basic and low-level. High level jobs should be done by running the separate programs or using separate functions in succession through a shell script or calling the libraries by higher level functions, see the examples in Tutorials. So when making the suggestions please consider how your desired job can best be broken into separate steps and modularized.

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1.9 Announcements

Gnuastro has a dedicated mailing list for making announcements (info-gnuastro). Anyone can subscribe to this mailing list. Anytime there is a new stable or test release, an email will be circulated there. The email contains a summary of the overall changes along with a detailed list (from the NEWS file). This mailing list is thus the best way to stay up to date with new releases, easily learn about the updated/new features, or dependencies (see Dependencies).

To subscribe to this list, please visit Traffic (number of mails per unit time) in this list is designed to be low: only a handful of mails per year. Previous announcements are available on its archive.

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1.10 Conventions

In this book we have the following conventions:

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1.11 Acknowledgments

The list of Gnuastro authors is available at the start of this book and the AUTHORS file in the source code. Here the authors wish to gratefully acknowledge the help and support they received from other people and institutions who had an indirect (not committed in the version controlled history) role in Gnuastro. The plain text file THANKS which is distributed along with the source code also contains this list.

The Japanese Ministry of Science and Technology (MEXT) scholarship for Mohammad Akhlaghi’s Masters and PhD period in Tohoku University Astronomical Institute had an instrumental role in the long term learning and planning that made the idea of Gnuastro possible. The very critical view points of Professor Takashi Ichikawa (Mohammad’s adviser) were also instrumental in the initial ideas and creation of Gnuastro. The European Research Council (ERC) advanced grant 339659-MUSICOS (Principal investigator: Roland Bacon) was vital in the growth and expansion of Gnuastro, enabling a thorough re-write of the core functionality of all libraries and programs, turning Gnuastro into the large collection of generic programs and libraries it is today.

In general, we would like to gratefully thank the following people for their useful and constructive comments and suggestions (in alphabetical order by family name): Valentina Abril-melgarejo, Marjan Akbari, Roland Bacon, Karl Berry, Leindert Boogaard, Nicolas Bouché, Fernando Buitrago, Adrian Bunk, Rosa Calvi, Nushkia Chamba, Benjamin Clement, Nima Dehdilani, Antonio Diaz Diaz, Thérèse Godefroy, Madusha Gunawardhana, Stephen Hamer, Takashi Ichikawa, Raúl Infante Sainz, Brandon Invergo, Oryna Ivashtenko, Aurélien Jarno, Lee Kelvin, Brandon Kelly, Mohammad-Reza Khellat, Geoffry Krouchi, Floriane Leclercq, Alan Lefor, Guillaume Mahler, Juan Molina Tobar, Francesco Montanari, Dmitrii Oparin, Bertrand Pain, William Pence, Bob Proulx, Teymoor Saifollahi, Yahya Sefidbakht, Alejandro Serrano Borlaff, Jenny Sorce, Lee Spitler, Richard Stallman, Ole Streicher, Alfred M. Szmidt, Michel Tallon, Juan C. Tello, Éric Thiébaut, Ignacio Trujillo, David Valls-Gabaud, Aaron Watkins, Christopher Willmer, Sara Yousefi Taemeh, Johannes Zabl. The GNU French Translation Team is also managing the French version of the top Gnuastro webpage which we highly appreciate. Finally we should thank all the (sometimes anonymous) people in various online forums which patiently answered all our small (but important) technical questions.

All work on Gnuastro has been voluntary, but the authors are most grateful to the following institutions (in chronological order) for hosting us in our research:

Ministry of education, culture, sports, science and technology (MEXT), Japan.
Tohoku University Astronomical Institute, Sendai, Japan.
University of Salento, Lecce, Italy.
Centre national de la recherche scientifique (CNRS), France.
Centre de Recherche Astrophysique de Lyon, University of Lyon 1, France.

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2 Tutorials

To help new users have a smooth and easy start with Gnuastro, in this chapter several thoroughly elaborated tutorials, or cookbooks, are provided. These tutorials demonstrate the capabilities of different Gnuastro programs and libraries, along with tips and guidelines for the best practices of using them in various realistic situations.

We strongly recommend going through these tutorials to get a good feeling of how the programs are related (built in a modular design to be used together in a pipeline), very similar to the core Unix-based programs that they were modeled on. Therefore these tutorials will greatly help in optimally using Gnuastro’s programs (and generally, the Unix-like command-line environment) effectively for your research.

In Sufi simulates a detection, we’ll start with a fictional21 tutorial explaining how Abd al-rahman Sufi (903 – 986 A.D., the first recorded description of “nebulous” objects in the heavens is attributed to him) could have used some of Gnuastro’s programs for a realistic simulation of his observations and see if his detection of nebulous objects was trust-able. Because all conditions are under control in a simulated/mock environment/dataset, mock datasets can be a valuable tool to inspect the limitations of your data analysis and processing. But they need to be as realistic as possible, so the first tutorial is dedicated to this important step of an analysis.

The next two tutorials (General program usage tutorial and Detecting large extended targets) use real input datasets from some of the deep Hubble Space Telescope (HST) images and the Sloan Digital Sky Survey (SDSS) respectively. Their aim is to demonstrate some real-world problems that many astronomers often face and how they can be be solved with Gnuastro’s programs.

The ultimate aim of General program usage tutorial is to detect galaxies in a deep HST image, measure their positions and brightness and select those with the strongest colors. In the process, it takes many detours to introduce you to the useful capabilities of many of the programs. So please be patient in reading it. If you don’t have much time and can only try one of the tutorials, we recommend this one.

Detecting large extended targets deals with a major problem in astronomy: effectively detecting the faint outer wings of bright (and large) nearby galaxies to extremely low surface brightness levels (roughly 1/20th of the local noise level in the example discussed). Besides the interesting scientific questions in these low-surface brightness features, failure to properly detect them will bias the measurements of the background objects and the survey’s noise estimates. This is an important issue, especially in wide surveys. Because bright/large galaxies and stars22, cover a significant fraction of the survey area.

Finally, in Hubble visually checks and classifies his catalog, we go into the historical/fictional world again to see how Hubble could have used Gnuastro’s programs to visually check and classify a sample of galaxies which ultimately lead him to the “Hubble fork” classification of galaxy morphologies.

In these tutorials, we have intentionally avoided too many cross references to make it more easy to read. For more information about a particular program, you can visit the section with the same name as the program in this book. Each program section in the subsequent chapters starts by explaining the general concepts behind what it does, for example see Convolve. If you only want practical information on running a program, for example its options/configuration, input(s) and output(s), please consult the subsection titled “Invoking ProgramName”, for example see Invoking NoiseChisel. For an explanation of the conventions we use in the example codes through the book, please see Conventions.

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2.1 Sufi simulates a detection

It is the year 953 A.D. and Sufi23 is in Shiraz as a guest astronomer. He had come there to use the advanced 123 centimeter astrolabe for his studies on the Ecliptic. However, something was bothering him for a long time. While mapping the constellations, there were several non-stellar objects that he had detected in the sky, one of them was in the Andromeda constellation. During a trip he had to Yemen, Sufi had seen another such object in the southern skies looking over the Indian ocean. He wasn’t sure if such cloud-like non-stellar objects (which he was the first to call ‘Sahābi’ in Arabic or ‘nebulous’) were real astronomical objects or if they were only the result of some bias in his observations. Could such diffuse objects actually be detected at all with his detection technique?

He still had a few hours left until nightfall (when he would continue his studies on the ecliptic) so he decided to find an answer to this question. He had thoroughly studied Claudius Ptolemy’s (90 – 168 A.D) Almagest and had made lots of corrections to it, in particular in measuring the brightness. Using his same experience, he was able to measure a magnitude for the objects and wanted to simulate his observation to see if a simulated object with the same brightness and size could be detected in a simulated noise with the same detection technique. The general outline of the steps he wants to take are:

  1. Make some mock profiles in an over-sampled image. The initial mock image has to be over-sampled prior to convolution or other forms of transformation in the image. Through his experiences, Sufi knew that this is because the image of heavenly bodies is actually transformed by the atmosphere or other sources outside the atmosphere (for example gravitational lenses) prior to being sampled on an image. Since that transformation occurs on a continuous grid, to best approximate it, he should do all the work on a finer pixel grid. In the end he can re-sample the result to the initially desired grid size.
  2. Convolve the image with a PSF image that is over-sampled to the same value as the mock image. Since he wants to finish in a reasonable time and the PSF kernel will be very large due to oversampling, he has to use frequency domain convolution which has the side effect of dimming the edges of the image. So in the first step above he also has to build the image to be larger by at least half the width of the PSF convolution kernel on each edge.
  3. With all the transformations complete, the image should be re-sampled to the same size of the pixels in his detector.
  4. He should remove those extra pixels on all edges to remove frequency domain convolution artifacts in the final product.
  5. He should add noise to the (until now, noise-less) mock image. After all, all observations have noise associated with them.

Fortunately Sufi had heard of GNU Astronomy Utilities from a colleague in Isfahan (where he worked) and had installed it on his computer a year before. It had tools to do all the steps above. He had used MakeProfiles before, but wasn’t sure which columns he had chosen in his user or system wide configuration files for which parameters, see Configuration files. So to start his simulation, Sufi runs MakeProfiles with the -P option to make sure what columns in a catalog MakeProfiles currently recognizes and the output image parameters. In particular, Sufi is interested in the recognized columns (shown below).

$ astmkprof -P

[[[ ... Truncated lines ... ]]]

# Output:
 type         float32     # Type of output: e.g., int16, float32, etc...
 naxis        1000,1000   # Number of pixels along first FITS axis.
 oversample   5           # Scale of oversampling (>0 and odd).

[[[ ... Truncated lines ... ]]]

# Columns, by info (see `--searchin'), or number (starting from 1):
 ccol         2           # Center along first FITS axis (horizontal).
 ccol         3           # Center along second FITS axis (vertical).
 fcol         4           # sersic (1), moffat (2), gaussian (3),
                          # point (4), flat (5), circumference (6).
 rcol         5           # Effective radius or FWHM in pixels.
 ncol         6           # Sersic index or Moffat beta.
 pcol         7           # Position angle.
 qcol         8           # Axis ratio.
 mcol         9           # Magnitude.
 tcol         10          # Truncation in units of radius or pixels.

[[[ ... Truncated lines ... ]]]

In Gnuastro, column counting starts from 1, so the columns are ordered such that the first column (number 1) can be an ID he specifies for each object (and MakeProfiles ignores), each subsequent column is used used for another property of the profile. It is also possible to use column names for the values of these options and change these defaults, but Sufi preferred to stick to the defaults. Fortunately MakeProfiles has the capability to also make the PSF which is to be used on the mock image and using the --prepforconv option, he can also make the mock image to be larger by the correct amount and all the sources to be shifted by the correct amount.

For his initial check he decides to simulate the nebula in the Andromeda constellation. The night he was observing, the PSF had roughly a FWHM of about 5 pixels, so as the first row (profile), he defines the PSF parameters and sets the radius column (rcol above, fifth column) to 5.000, he also chooses a Moffat function for its functional form. Remembering how diffuse the nebula in the Andromeda constellation was, he decides to simulate it with a mock Sérsic index 1.0 profile. He wants the output to be 500 pixels by 500 pixels, so he puts the mock profile in the center. Looking at his drawings of it, he decides a reasonable effective radius for it would be 40 pixels on this image pixel scale, he sets the axis ratio and position angle to approximately correct values too and finally he sets the total magnitude of the profile to 3.44 which he had accurately measured. Sufi also decides to truncate both the mock profile and PSF at 5 times the respective radius parameters. In the end he decides to put four stars on the four corners of the image at very low magnitudes as a visual scale.

Using all the information above, he creates the catalog of mock profiles he wants in a file named cat.txt (short for catalog) using his favorite text editor and stores it in a directory named simulationtest in his home directory. [The cat command prints the contents of a file, short for concatenation. So please copy-paste the lines after “cat cat.txt” into cat.txt when the editor opens in the steps above it, note that there are 7 lines, first one starting with #]:

$ mkdir ~/simulationtest
$ cd ~/simulationtest
$ pwd
$ emacs cat.txt
$ ls
$ cat cat.txt
# Column 4: PROFILE_NAME [,str7] Radial profile's functional name
 1  0.0000   0.0000  moffat  5.000  4.765  0.0000  1.000  30.000  5.000
 2  250.00   250.00  sersic  40.00  1.000  -25.00  0.400  3.4400  5.000
 3  50.000   50.000  point   0.000  0.000  0.0000  0.000  9.0000  0.000
 4  450.00   50.000  point   0.000  0.000  0.0000  0.000  9.2500  0.000
 5  50.000   450.00  point   0.000  0.000  0.0000  0.000  9.5000  0.000
 6  450.00   450.00  point   0.000  0.000  0.0000  0.000  9.7500  0.000

The zero-point magnitude for his observation was 18. Now he has all the necessary parameters and runs MakeProfiles with the following command:

$ astmkprof --prepforconv --naxis=500,500 --zeropoint=18.0 cat.txt
MakeProfiles started on Sat Oct  6 16:26:56 953
  - 6 profiles read from cat.txt
  - Random number generator (RNG) type: mt19937
  - Using 8 threads.
  ---- row 2 complete, 5 left to go
  ---- row 3 complete, 4 left to go
  ---- row 4 complete, 3 left to go
  ---- row 5 complete, 2 left to go
  ---- ./0_cat.fits created.
  ---- row 0 complete, 1 left to go
  ---- row 1 complete, 0 left to go
  - ./cat.fits created.                                0.041651 seconds
MakeProfiles finished in 0.267234 seconds

0_cat.fits  cat.fits  cat.txt

The file 0_cat.fits is the PSF Sufi had asked for and cat.fits is the image containing the other 5 objects. The PSF is now available to him as a separate file for the convolution step. While he was preparing the catalog, one of his students approached him and was also following the steps. When he opened the image, the student was surprised to see that all the stars are only one pixel and not in the shape of the PSF as we see when we image the sky at night. So Sufi explained to him that the stars will take the shape of the PSF after convolution and this is how they would look if we didn’t have an atmosphere or an aperture when we took the image. The size of the image was also surprising for the student, instead of 500 by 500, it was 2630 by 2630 pixels. So Sufi had to explain why oversampling is important for parts of the image where the flux change is significant over a pixel. Sufi then explained to him that after convolving we will re-sample the image to get our originally desired size. To convolve the image, Sufi ran the following command:

$ astconvolve --kernel=0_cat.fits cat.fits
Convolve started on Mon Apr  6 16:35:32 953
  - Using 8 CPU threads.
  - Input: cat.fits (hdu: 1)
  - Kernel: 0_cat.fits (hdu: 1)
  - Input and Kernel images padded.                    0.075541 seconds
  - Images converted to frequency domain.              6.728407 seconds
  - Multiplied in the frequency domain.                0.040659 seconds
  - Converted back to the spatial domain.              3.465344 seconds
  - Padded parts removed.                              0.016767 seconds
Convolve finished in:  10.422161 seconds

0_cat.fits  cat_convolved.fits  cat.fits  cat.txt

When convolution finished, Sufi opened the cat_convolved.fits file and showed the effect of convolution to his student and explained to him how a PSF with a larger FWHM would make the points even wider. With the convolved image ready, they were prepared to re-sample it to the original pixel scale Sufi had planned [from the $ astmkprof -P command above, recall that MakeProfiles had over-sampled the image by 5 times]. Sufi explained the basic concepts of warping the image to his student and ran Warp with the following command:

$ astwarp --scale=1/5 --centeroncorner cat_convolved.fits
Warp started on Mon Apr  6 16:51:59 953
 Using 8 CPU threads.
 Input: cat_convolved.fits (hdu: 1)
        0.2000   0.0000   0.4000
        0.0000   0.2000   0.4000
        0.0000   0.0000   1.0000

$ ls
0_cat.fits          cat_convolved_scaled.fits     cat.txt
cat_convolved.fits  cat.fits

$ astfits -p cat_convolved_scaled.fits | grep NAXIS
NAXIS   =                    2 / number of data axes
NAXIS1  =                  526 / length of data axis 1
NAXIS2  =                  526 / length of data axis 2

cat_convolved_warped.fits now has the correct pixel scale. However, the image is still larger than what we had wanted, it is 526 (\(500+13+13\)) by 526 pixels. The student is slightly confused, so Sufi also re-samples the PSF with the same scale and shows him that it is 27 (\(2\times13+1\)) by 27 pixels. Sufi goes on to explain how frequency space convolution will dim the edges and that is why he added the --prepforconv option to MakeProfiles, see If convolving afterwards. Now that convolution is done, Sufi can remove those extra pixels using Crop with the command below. Crop’s --section option accepts coordinates inclusively and counting from 1 (according to the FITS standard), so the crop’s first pixel has to be 14, not 13.

$ astcrop cat_convolved_scaled.fits --section=14:*-13,14:*-13    \
Crop started on Sat Oct  6 17:03:24 953
  - Read metadata of 1 image.                          0.001304 seconds
  ---- ...nvolved_scaled_cropped.fits created: 1 input.
Crop finished in:  0.027204 seconds

0_cat.fits          cat_convolved_scaled_cropped.fits  cat.fits
cat_convolved.fits  cat_convolved_scaled.fits          cat.txt

Finally, cat_convolved_scaled_cropped.fits has the same dimensions as Sufi had desired in the beginning. All this trouble was certainly worth it because now there is no dimming on the edges of the image and the profile centers are more accurately sampled. The final step to simulate a real observation would be to add noise to the image. Sufi set the zeropoint magnitude to the same value that he set when making the mock profiles and looking again at his observation log, he had measured the background flux near the nebula had a magnitude of 7 that night. So using these values he ran MakeNoise:

$ astmknoise --zeropoint=18 --background=7 --output=out.fits    \
MakeNoise started on Mon Apr  6 17:05:06 953
  - Generator type: mt19937
  - Generator seed: 1428318100
MakeNoise finished in:  0.033491 (seconds)

0_cat.fits         cat_convolved_scaled_cropped.fits cat.fits  out.fits
cat_convolved.fits cat_convolved_scaled.fits         cat.txt

The out.fits file now contains the noised image of the mock catalog Sufi had asked for. Seeing how the --output option allows the user to specify the name of the output file, the student was confused and wanted to know why Sufi hadn’t used it before? Sufi then explained to him that for intermediate steps it is best to rely on the automatic output, see Automatic output. Doing so will give all the intermediate files the same basic name structure, so in the end you can simply remove them all with the Shell’s capabilities. So Sufi decided to show this to the student by making a shell script from the commands he had used before.

The command-line shell has the capability to read all the separate input commands from a file. This is useful when you want to do the same thing multiple times, with only the names of the files or minor parameters changing between the different instances. Using the shell’s history (by pressing the up keyboard key) Sufi reviewed all the commands and then he retrieved the last 5 commands with the $ history 5 command. He selected all those lines he had input and put them in a text file named Then he defined the edge and base shell variables for easier customization later. Finally, before every command, he added some comments (lines starting with #) for future readability.

# Basic settings:

# Remove any existing image to avoid confusion.
rm out.fits

# Run MakeProfiles to create an oversampled FITS image.
astmkprof --prepforconv --naxis=500,500 --zeropoint=18.0 "$base".txt

# Convolve the created image with the kernel.
astconvolve --kernel=0_"$base".fits "$base".fits

# Scale the image back to the intended resolution.
astwarp --scale=1/5 --centeroncorner "$base"_convolved.fits

# Crop the edges out (dimmed during convolution). `--section' accepts
# inclusive coordinates, so the start of start of the section must be
# one pixel larger than its end.
st_edge=$(( edge + 1 ))
astcrop "$base"_convolved_scaled.fits --zeroisnotblank          \

# Add noise to the image.
astmknoise --zeropoint=18 --background=7 --output=out.fits      \

# Remove all the temporary files.
rm 0*.fits cat*.fits

He used this chance to remind the student of the importance of comments in code or shell scripts: when writing the code, you have a good mental picture of what you are doing, so writing comments might seem superfluous and excessive. However, in one month when you want to re-use the script, you have lost that mental picture and rebuilding it is can be time-consuming and frustrating. The importance of comments is further amplified when you want to share the script with a friend/colleague. So it is good to accompany any script/code with useful comments while you are writing it (have a good mental picture of what/why you are doing something).

Sufi then explained to the eager student that you define a variable by giving it a name, followed by an = sign and the value you want. Then you can reference that variable from anywhere in the script by calling its name with a $ prefix. So in the script whenever you see $base, the value we defined for it above is used. If you use advanced editors like GNU Emacs or even simpler ones like Gedit (part of the GNOME graphical user interface) the variables will become a different color which can really help in understanding the script. We have put all the $base variables in double quotation marks (") so the variable name and the following text do not get mixed, the shell is going to ignore the " after replacing the variable value. To make the script executable, Sufi ran the following command:

$ chmod +x

Then finally, Sufi ran the script, simply by calling its file name:

$ ./

After the script finished, the only file remaining is the out.fits file that Sufi had wanted in the beginning. Sufi then explained to the student how he could run this script anywhere that he has a catalog if the script is in the same directory. The only thing the student had to modify in the script was the name of the catalog (the value of the base variable in the start of the script) and the value to the edge variable if he changed the PSF size. The student was also happy to hear that he won’t need to make it executable again when he makes changes later, it will remain executable unless he explicitly changes the executable flag with chmod.

The student was really excited, since now, through simple shell scripting, he could really speed up his work and run any command in any fashion he likes allowing him to be much more creative in his works. Until now he was using the graphical user interface which doesn’t have such a facility and doing repetitive things on it was really frustrating and some times he would make mistakes. So he left to go and try scripting on his own computer.

Sufi could now get back to his own work and see if the simulated nebula which resembled the one in the Andromeda constellation could be detected or not. Although it was extremely faint24, fortunately it passed his detection tests and he wrote it in the draft manuscript that would later become “Book of fixed stars”. He still had to check the other nebula he saw from Yemen and several other such objects, but they could wait until tomorrow (thanks to the shell script, he only has to define a new catalog). It was nearly sunset and they had to begin preparing for the night’s measurements on the ecliptic.

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2.2 General program usage tutorial

Measuring colors of astronomical objects in broad-band or narrow-band images is one of the most basic and common steps in astronomical analysis. Here, we will use Gnuastro’s programs to get a physical scale (area at certain redshifts) of the field we are studying, detect objects in a Hubble Space Telescope (HST) image, measure their colors and identify the ones with the largest colors to visual inspection and their spatial position in the image. After this tutorial, you can also try the Detecting large extended targets tutorial which goes into a little more detail on optimally configuring NoiseChisel (Gnuastro’s detection tool) in special situations.

During the tutorial, we will take many detours to explain, and practically demonstrate, the many capabilities of Gnuastro’s programs. In the end you will see that the things you learned during this tutorial are much more generic than this particular problem and can be used in solving a wide variety of problems involving the analysis of data (images or tables). So please don’t rush, and go through the steps patiently to optimally master Gnuastro.

In this tutorial, we’ll use the HST eXtreme Deep Field dataset. Like almost all astronomical surveys, this dataset is free for download and usable by the public. You will need the following tools in this tutorial: Gnuastro, SAO DS9 25, GNU Wget26, and AWK (most common implementation is GNU AWK27).

This tutorial was first prepared for the “Exploring the Ultra-Low Surface Brightness Universe” workshop (November 2017) at the ISSI in Bern, Switzerland. It was further extended in the “4th Indo-French Astronomy School” (July 2018) organized by LIO, CRAL CNRS UMR5574, UCBL, and IUCAA in Lyon, France. We are very grateful to the organizers of these workshops and the attendees for the very fruitful discussions and suggestions that made this tutorial possible.

Write the example commands manually: Try to type the example commands on your terminal manually and use the history feature of your command-line (by pressing the “up” button to retrieve previous commands). Don’t simply copy and paste the commands shown here. This will help simulate future situations when you are processing your own datasets.

A handy feature of Gnuastro is that all program names start with ast. This will allow your command-line processor to easily list and auto-complete Gnuastro’s programs for you. Try typing the following command (press TAB key when you see <TAB>) to see the list:

$ ast<TAB><TAB>

Any program that starts with ast (including all Gnuastro programs) will be shown. By choosing the subsequent characters of your desired program and pressing <TAB><TAB> again, the list will narrow down and the program name will auto-complete once your input characters are unambiguous. In short, you often don’t need to type the full name of the program you want to run.

Gnuastro contains a large number of programs and it is natural to forget the details of each program’s options or inputs and outputs. Therefore, before starting the analysis, let’s review how you can access this book to refresh your memory any time you want. For example when working on the command-line, without having to take your hands off the keyboard. When you install Gnuastro, this book is also installed on your system along with all the programs and libraries, so you don’t need an internet connection to to access/read it. Also, by accessing this book as described below, you can be sure that it corresponds to your installed version of Gnuastro.

GNU Info28 is the program in charge of displaying the manual on the command-line (for more, see Info). To see this whole book on your command-line, please run the following command and press subsequent keys. Info has its own mini-environment, therefore we’ll show the keys that must be pressed in the mini-environment after a -> sign. You can also ignore anything after the # sign in the middle of the line, they are only for your information.

$ info gnuastro                # Open the top of the manual.
-> <SPACE>                     # All the book chapters.
-> <SPACE>                     # Continue down: show sections.
-> <SPACE> ...                 # Keep pressing space to go down.
-> q                           # Quit Info, return to the command-line.

The thing that greatly simplifies navigation in Info is the links (regions with an underline). You can immediately go to the next link in the page with the <TAB> key and press <ENTER> on it to go into that part of the manual. Try the commands above again, but this time also use <TAB> to go to the links and press <ENTER> on them to go to the respective section of the book. Then follow a few more links and go deeper into the book. To return to the previous page, press l (small L). If you are searching for a specific phrase in the whole book (for example an option name), press s and type your search phrase and end it with an <ENTER>.

You don’t need to start from the top of the manual every time. For example, to get to Invoking NoiseChisel, run the following command. In general, all programs have such an “Invoking ProgramName” section in this book. These sections are specifically for the description of inputs, outputs and configuration options of each program. You can access them directly for each program by giving its executable name to Info.

$ info astnoisechisel

The other sections don’t have such shortcuts. To directly access them from the command-line, you need to tell Info to look into Gnuastro’s manual, then look for the specific section (an unambiguous title is necessary). For example, if you only want to review/remember NoiseChisel’s Detection options), just run the following command. Note how case is irrelevant for Info when calling a title in this manner.

$ info gnuastro "Detection options"

In general, Info is a powerful and convenient way to access this whole book with detailed information about the programs you are running. If you are not already familiar with it, please run the following command and just read along and do what it says to learn it. Don’t stop until you feel sufficiently fluent in it. Please invest the half an hour’s time necessary to start using Info comfortably. It will greatly improve your productivity and you will start reaping the rewards of this investment very soon.

$ info info

As a good scientist you need to feel comfortable to play with the features/options and avoid (be critical to) using default values as much as possible. On the other hand, our human memory is limited, so it is important to be able to easily access any part of this book fast and remember the option names, what they do and their acceptable values.

If you just want the option names and a short description, calling the program with the --help option might also be a good solution like the first example below. If you know a few characters of the option name, you can feed the output to grep like the second or third example commands.

$ astnoisechisel --help
$ astnoisechisel --help | grep quant
$ astnoisechisel --help | grep check

Let’s start the processing. First, to keep things clean, let’s create a gnuastro-tutorial directory and continue all future steps in it:

$ mkdir gnuastro-tutorial
$ cd gnuastro-tutorial

We will be using the near infra-red Wide Field Camera dataset. If you already have them in another directory (for example XDFDIR), you can set the download directory to be a symbolic link to XDFDIR with a command like this:

$ ln -s XDFDIR download

If the following images aren’t already present on your system, you can make a download directory and download them there.

$ mkdir download
$ cd download
$ xdfurl=
$ wget $xdfurl/hlsp_xdf_hst_wfc3ir-60mas_hudf_f105w_v1_sci.fits
$ wget $xdfurl/hlsp_xdf_hst_wfc3ir-60mas_hudf_f160w_v1_sci.fits
$ cd ..

In this tutorial, we’ll just use these two filters. Later, you will probably need to download more filters, you can use the shell’s for loop to download them all in series (one after the other29) with one command like the one below for the WFC3 filters. Put this command instead of the two wget commands above. Recall that all the extra spaces, back-slashes (\), and new lines can be ignored if you are typing on the lines on the terminal.

$ for f in f105w f125w f140w f160w; do                              \
    wget $xdfurl/hlsp_xdf_hst_wfc3ir-60mas_hudf_"$f"_v1_sci.fits;   \

First, let’s visually inspect the dataset. Let’s take F160W image as an example. Do the steps below with the other image(s) too (and later with any dataset that you want to work on). It is very important to understand your dataset visually. Note how ds9 doesn’t follow the GNU style of options where “long” and “short” options are preceded by -- and - respectively (for example --width and -w, see Options).

Ds9’s -zscale option is a good scaling to highlight the low surface brightness regions, and as the name suggests, -zoom to fit will fit the whole dataset in the window. If the window is too small, expand it with your mouse, then press the “zoom” button on the top row of buttons above the image, then in the row below it, press “zoom fit”. You can also zoom in and out by scrolling your mouse or the respective operation on your touch-pad when your cursor/pointer is over the image.

$ ds9 download/hlsp_xdf_hst_wfc3ir-60mas_hudf_f160w_v1_sci.fits     \
      -zscale -zoom to fit

The first thing you might notice is that the regions with no data have a value of zero in this image. The next thing might be that the dataset actually has two “depth”s (see Quantifying measurement limits). The exposure time of the inner region is more than 4 times of the outer parts. Fortunately the XDF survey webpage (above) contains the vertices of the deep flat WFC3-IR field. You can use those vertices in Crop to cutout this deep infra-red region from the larger image. We’ll make a directory called flat-ir and keep the flat infra-red regions in that directory (with a ‘xdf-’ suffix for a shorter and easier filename).

$ mkdir flat-ir
$ astcrop --mode=wcs -h0 --output=flat-ir/xdf-f105w.fits              \
          --polygon="53.187414,-27.779152 : 53.159507,-27.759633 :    \
                     53.134517,-27.787144 : 53.161906,-27.807208"     \
$ astcrop --mode=wcs -h0 --output=flat-ir/xdf-f160w.fits              \
          --polygon="53.187414,-27.779152 : 53.159507,-27.759633 :    \
                     53.134517,-27.787144 : 53.161906,-27.807208"     \

The only thing varying in the two calls to Gnuastro’s Crop program is the filter name. Therefore, to simplify the command, and later allow work on more filters, we can use the shell’s for loop. Notice how the two places where the filter names (f105w and f160w) are used above have been replaced with $f (the shell variable that for is in charge of setting) below. To generalize this for more filters later, you can simply add the other filter names in the first line before the semi-colon (;).

$ for f in f105w f160w; do                                            \
    astcrop --mode=wcs -h0 --output=flat-ir/xdf-$f.fits               \
            --polygon="53.187414,-27.779152 : 53.159507,-27.759633 :  \
                       53.134517,-27.787144 : 53.161906,-27.807208"   \
            download/hlsp_xdf_hst_wfc3ir-60mas_hudf_"$f"_v1_sci.fits; \

Please open these images and inspect them with the same ds9 commands you used above. You will see how it is completely flat now and doesn’t have varying depths. Another important result of this crop is that regions with no data now have a NaN (blank) value, not zero. Zero is a meaningful value and especially when using NoiseChisel, the input should have NaN values for pixels with no data, not zero.

This is the deepest image we currently have of the sky. The first thing that comes to mind may be this: “How large is this field?”. Let’s find the answer to this question with the commands below. The lines starting with ## are just comments for you to help in following the steps. Don’t type them on the terminal. The commands are intentionally repetitive in some places to better understand each step and also to demonstrate the beauty of command-line features like variables, pipes and loops. Later, if you would like to repeat this process on another dataset, you can just use commands 3, 7, and 9.

Use shell history: Don’t forget to make effective use of your shell’s history. This is especially convenient when you just want to make a small change to your previous command. Press the “up” key on your keyboard (possibly multiple times) to see your previous command(s).

## (1)  See the general statistics of non-blank pixel values.
$ aststatistics flat-ir/xdf-f160w.fits

## (2)  We only want the number of non-blank pixels.
$ aststatistics flat-ir/xdf-f160w.fits --number

## (3)  Keep the result of the command above in the shell variable `n'.
$ n=$(aststatistics flat-ir/xdf-f160w.fits --number)

## (4)  See what is stored the shell variable `n'.
$ echo $n

## (5)  Show all the FITS keywords of this image.
$ astfits flat-ir/xdf-f160w.fits -h1

## (6)  The resolution (in degrees/pixel) is in the `CDELT' keywords.
##      Only show lines that contain these characters, by feeding
##      the output of the previous command to the `grep' program.
$ astfits flat-ir/xdf-f160w.fits -h1 | grep CDELT

## (7)  Save the resolution (same in both dimensions) in the variable
##      `r'. The last part uses AWK to print the third `field' of its
##      input line. The first two fields were `CDELT1' and `='.
$ r=$(astfits flat-ir/xdf-f160w.fits -h1 | grep CDELT1   \
              | awk '{print $3}')

## (8)  Print the values of `n' and `r'.
$ echo $n $r

## (9)  Use the number of pixels (first number passed to AWK) and
##      length of each pixel's edge (second number passed to AWK)
##      to estimate the area of the field in arc-minutes squared.
$ area=$(echo $n $r | awk '{print $1 * ($2^2) * 3600}')

The area of this field is 4.03817 (or 4.04) arc-minutes squared. Just for comparison, this is roughly 175 times smaller than the average moon’s angular area (with a diameter of 30arc-minutes or half a degree).

AWK for table/value processing: AWK is a powerful and simple tool for text processing. Above (and further below) some simple examples are shown. GNU AWK (the most common implementation) comes with a free and wonderful book in the same format as this book which will allow you to master it nicely. Just like this manual, you can also access GNU AWK’s manual on the command-line whenever necessary without taking your hands off the keyboard.

This takes us to the second question that you have probably asked yourself when you saw the field for the first time: “How large is this area at different redshifts?”. To get a feeling of the tangential area that this field covers at redshift 2, you can use CosmicCalculator. In particular, you need the tangential distance covered by 1 arc-second as raw output. Combined with the field’s area, we can then calculate the tangential distance in Mega Parsecs squared (\(Mpc^2\)).

## Print general cosmological properties at redshift 2.
$ astcosmiccal -z2

## When given a "Specific calculation" option, CosmicCalculator
## will just print that particular calculation. See the options
## under this title in the output of `--help' for more.
$ astcosmiccal --help

## Only print the "Tangential dist. covered by 1arcsec at z (kpc)".
## in units of kpc/arc-seconds.
$ astcosmiccal -z2 --arcsectandist

## Convert this distance to kpc^2/arcmin^2 and save in `k'.
$ k=$(astcosmiccal -z2 --arcsectandist | awk '{print ($1*60)^2}')

## Multiply by the area of the field (in arcmin^2) and divide by
## 10^6 to return value in Mpc^2.
$ echo $k $area | awk '{print $1 * $2 / 1e6}'

At redshift 2, this field therefore covers 1.07145 \(Mpc^2\). If you would like to see how this tangential area changes with redshift, you can use a shell loop like below.

$ for z in 0.5 1.0 1.5 2.0 2.5 3.0 3.5 4.0 4.5 5.0; do           \
    k=$(astcosmiccal -z$z --arcsectandist);                      \
    echo $z $k $area | awk '{print $1, ($2*60)^2 * $3 / 1e6}';   \

Fortunately, the shell has a useful tool/program to print a sequence of numbers that is nicely called seq. You can use it instead of typing all the different redshifts in this example. For example the loop below will print the same range of redshifts (between 0.5 and 5) but with increments of 0.1.

$ for z in $(seq 0.5 0.1 5); do                                  \
    k=$(astcosmiccal -z$z --arcsectandist);                      \
    echo $z $k $area | awk '{print $1, ($2*60)^2 * $3 / 1e6}';   \

This is a fast and simple way for this repeated calculation when it is only necessary once. However, if you commonly need this calculation and possibly for a larger number of redshifts, the command above can be slow. This is because the CosmicCalculator program has a lot of overhead. To be generic and easy to operate, it has to parse the command-line and all configuration files (see below) which contain human-readable characters and need a lot of processing to be ready for processing by the computer. Afterwards, CosmicCalculator has to check the sanity of its inputs and check which of its many options you have asked for. It has to do all of these for every redshift in the loop above.

To greatly speed up the processing, you can directly access the root work-horse of CosmicCalculator without all that overhead. Using Gnuastro’s library, you can write your own tiny program particularly designed for this exact calculation (and nothing else!). To do that, copy and paste the following C program in a file called myprogram.c.

#include <math.h>
#include <stdio.h>
#include <stdlib.h>
#include <gnuastro/cosmology.h>

  double area=4.03817;          /* Area of field (arcmin^2). */
  double z, adist, tandist;     /* Temporary variables.      */

  /* Constants from Plank 2018 (arXiv:1807.06209, Table 2) */
  double H0=67.66, olambda=0.6889, omatter=0.3111, oradiation=0;

  /* Do the same thing for all redshifts (z) between 0.1 and 5. */
  for(z=0.1; z<5; z+=0.1)
      /* Calculate the angular diameter distance. */
      adist=gal_cosmology_angular_distance(z, H0, olambda,
                                           omatter, oradiation);

      /* Calculate the tangential distance of one arcsecond. */
      tandist = adist * 1000 * M_PI / 3600 / 180;

      /* Print the redshift and area. */
      printf("%-5.2f %g\n", z, pow(tandist * 60,2) * area / 1e6);

  /* Tell the system that everything finished successfully. */
  return EXIT_SUCCESS;

To greatly simplify the compilation, linking and running of simple C programs like this that use Gnuastro’s library, Gnuastro has BuildProgram. This program designed to manage Gnuastro’s dependencies, compile and link the program and then run the new program. To build and run the program above, simply run the following command:

$ astbuildprog myprogram.c

Did you notice how much faster this was compared to the shell loop we wrote above? You might have noticed that a new file called myprogram is also created in the directory. This is the compiled program that was created and run by the command above (its in binary machine code format, not human-readable any more). You can run it again to get the same results with a command like this:

$ ./myprogram

The efficiency of myprogram compared to CosmicCalculator is because the requested processing is faster/comparable to the overheads necessary for each processing. For other programs that take large input datasets (images for example), the overhead is usually negligible compared to the processing. In such cases, the libraries are only useful if you want a different/new processing compared to the functionalities in Gnuastro’s existing programs.

Gnuastro has a large library which is heavily used by all the programs. In other words, the library is like the skeleton of Gnuastro. For the full list of available functions classified by context, please see Gnuastro library. Gnuastro’s library and BuildProgram are created to make it easy for you to use these powerful features as you like. This gives you a high level of creativity, while also providing efficiency and robustness. Several other complete working examples (involving images and tables) of Gnuastro’s libraries can be see in Library demo programs. Let’s stop the discussion on libraries at this point in this tutorial and get back to Gnuastro’s already built programs which were the main purpose of this tutorial.

None of Gnuastro’s programs keep a default value internally within their code. However, when you ran CosmicCalculator with the -z2 option above, it completed its processing and printed results. So where did the “default” cosmological parameter values (like the matter density and etc) come from? The values come from the command-line or a configuration file (see Configuration file precedence).

CosmicCalculator has a limited set of parameters and doesn’t need any particular file inputs. Therefore, we’ll use it to discuss configuration files which are an important part of all Gnuastro’s programs (see Configuration files).

Once you get comfortable with configuration files, you can easily do the same for the options of all Gnuastro programs. For example, NoiseChisel has the largest number of options in the programs. Therefore configuration files will be useful for it when you use different datasets (with different noise properties or in different research contexts). The configuration of each program (besides its version) is vital for the reproducibility of your results, so it is important to manage them properly.

As we saw above, the full list of the options in all Gnuastro programs can be seen with the --help option. Try calling it with CosmicCalculator as shown below. Note how options are grouped by context to make it easier to find your desired option. However, in each group, options are ordered alphabetically.

$ astcosmiccal --help

The options that need a value have an = sign after their long version and FLT, INT or STR for floating point numbers, integer numbers and strings (filenames for example) respectively. All options have a long format and some have a short format (a single character), for more see Options.

When you are using a program, it is often necessary to check the value the option has just before the program starts its processing. In other words, after it has parsed the command-line options and all configuration files. You can see the values of all options that need one with the --printparams or -P option that is common to all programs (see Common options). In the command below, try replacing -P with --printparams to see how both do the same operation.

$ astcosmiccal -P

Let’s say you want a different Hubble constant. Try running the following command to see how the Hubble constant in the output of the command above has changed. Afterwards, delete the -P and add a -z2 to see the results with the new cosmology (or configuration).

$ astcosmiccal -P --H0=70

From the output of the --help option, note how the option for Hubble constant has both short (-H) and long (--H0) formats. One final note is that the equal (=) sign is not mandatory. In the short format, the value can stick to the actual option (the short option name is just one character after-all and thus easily identifiable) and in the long format, a white-space character is also enough.

$ astcosmiccal -H70    -z2
$ astcosmiccal --H0 70 -z2 --arcsectandist

Let’s assume that in one project, you want to only use rounded cosmological parameters (H0 of 70km/s/Mpc and matter density of 0.3). You should therefore run CosmicCalculator like this:

$ astcosmiccal --H0=70 --olambda=0.7 --omatter=0.3 -z2

But having to type these extra options every time you run CosmicCalculator will be prone to errors (typos in particular) and also will be frustrating and slow. Therefore in Gnuastro, you can put all the options and their values in a “Configuration file” and tell the programs to read the option values from there.

Let’s create a configuration file. In your favorite text editor, make a file named my-cosmology.conf (or my-cosmology.txt, the suffix doesn’t matter) which contains the following lines. One space between the option value and name is enough, the values are just under each other to help in readability. Also note that you can only use long option names in configuration files.

H0       70
olambda  0.7
omatter  0.3

You can now tell CosmicCalculator to read this file for option values immediately using the --config option as shown below. Do you see how the output of the following command corresponds to the option values in my-cosmology.txt (previous command)?

$ astcosmiccal --config=my-cosmology.conf -z2

If you need this cosmology every time you are working in a specific directory, you can benefit from Gnuastro’s default configuration files to avoid having to call the --config option. Let’s assume that you want any CosmicCalculator call you make in the my-cosmology directory to use these parameters. You just have to copy the above configuration file into a special directory and file:

$ mkdir my-cosmology
$ mkdir my-cosmology/.gnuastro
$ mv my-cosmology.conf my-cosmology/.gnuastro/astcosmiccal.conf

Once you run CosmicCalculator within my-cosmology as shown below, you will see how your cosmology has been implemented without having to type anything extra on the command-line.

$ cd my-cosmology
$ astcosmiccal -z2
$ cd ..

To further simplify the process, you can use the --setdirconf option. If you are already in your desired directory, calling this option with the others will automatically write the final values (along with descriptions) in .gnuastro/astcosmiccal.conf. For example the commands below will make the same configuration file automatically (with one extra call to CosmicCalculator).

$ mkdir my-cosmology2
$ cd my-cosmology2
$ astcosmiccal --H0 70 --olambda=0.7 --omatter=0.3 --setdirconf
$ astcosmiccal -z2
$ cd ..

Gnuastro’s programs also have default configuration files for a specific user (when run in any directory). This allows you to set a special behavior every time a program is run by a specific user. Only the directory and filename differ from the above, the rest of the process is similar to before. Finally, there are also system-wide configuration files that can be used to define the option values for all users on a system. See Configuration file precedence for a more detailed discussion.

We are now ready to start processing the downloaded images. Since these datasets are already aligned, you don’t need to align them to make sure the pixel grid covers the same region in all inputs. Gnuastro’s Warp program has features for such pixel-grid warping (see Warp). Therefore, just for a demonstration, let’s assume one image needs to be rotated by 20 degrees to correspond to the other. To do that, you can run the following command:

$ astwarp flat-ir/xdf-f160w.fits --rotate=20

Open the output and see the output. If your final image is already aligned with RA and Dec, you can simply use the --align option and let Warp calculate the necessary rotation.

Warp can generally be used for any kind of pixel grid manipulation (warping). For example the outputs of the commands below will respectively have larger pixels (new resolution being one quarter the original resolution), get shifted by 2.8 (by sub-pixel), get a shear of 2, and be tilted (projected). After running each, please open the output file and see the effect.

$ astwarp flat-ir/xdf-f160w.fits --scale=0.25
$ astwarp flat-ir/xdf-f160w.fits --translate=2.8
$ astwarp flat-ir/xdf-f160w.fits --shear=2
$ astwarp flat-ir/xdf-f160w.fits --project=0.001,0.0005

You can also combine multiple warps in one command. For example to first rotate the image, then scale it, run this command:

$ astwarp flat-ir/xdf-f160w.fits --rotate=20 --scale=0.25

If you have multiple warps, do them all in one command. Don’t warp them in separate commands because the correlated noise will become too strong. As you see in the matrix that is printed when you run Warp, it merges all the warps into a single warping matrix (see Warping basics and Merging multiple warpings) and simply applies that just once. Recall that since this is done through matrix multiplication, order matters in the separate operations. In fact through Warp’s --matrix option, you can directly request your desired final warp and don’t have to break it up into different warps like above (see Invoking Warp).

Fortunately these datasets are already aligned to the same pixel grid, so you don’t actually need the files that were just generated. You can safely delete them all with the following command. Here, you see why we put the processed outputs that we need later into a separate directory. In this way, the top directory can be used for temporary files for testing that you can simply delete with a generic command like below.

$ rm *.fits

To detect the signal in the image (separate interesting pixels from noise), we’ll run NoiseChisel (NoiseChisel):

$ astnoisechisel flat-ir/xdf-f160w.fits --output=nc-test.fits

NoiseChisel’s output is a single file containing multiple extensions. You can get basic information about the extensions in a FITS file with Gnuastro’s Fits program (see Fits):

$ astfits nc-test.fits

From the output list, you see NoiseChisel’s output contains 5 extensions and the first (counted as zero) is blank: it has value of 0 in the last/size column, showing that it contains no data, it just contains meta-data. All of Gnuastro’s programs follow this convention of writing no data in the first HDU/extension. This allows the first extension to keep meta-data about all the extensions and is recommended by the FITS standard, see Fits.

The name of each extension describes its contents: the first (INPUT-NO-SKY) is the Sky-subtracted input that you provided. The second (NoiseChisel’s main output, DETECTIONS) has a numeric data type of uint8 with only two possible values for all pixels: 0 for noise and 1 for signal. The third and fourth (called SKY and SKY_STD), have the Sky and its standard deviation values of the input on a tessellation and were calculated over the undetected regions.

To better understand NoiseChisel’s output and the extensions described above, let’s visually inspect it (here, we’ll use SAO DS9). Since the file contains multiple related extensions, the easiest way to view all of them in DS9 is to open it as a “Multi-extension data cube” with the -mecube option as shown below30.

$ ds9 -mecube nc-test.fits -zscale -zoom to fit

The buttons and horizontal scroll bar in the “cube” window can be used to navigate between the extensions. In this mode, all DS9’s settings (for example zoom or color-bar) will be identical between the extensions. Try zooming into to one part and seeing how the galaxies were detected along with the Sky and Sky standard deviation values. Just have in mind that NoiseChisel’s job is only detection (separating signal from noise), We’ll segment this result later.

One good way to see if you have missed any signal is to mask all the detected pixels and inspect the noise pixels. For this, you can use Gnuastro’s Arithmetic program (in particular its where operator, see Arithmetic operators). With the command below, all detected pixels (in the DETECTIONS extension) will be set to NaN in the output (nc-masked.fits). To make the command easier to read/write, let’s just put the file name in a shell variable (img) first. A shell variable’s value can be retrieved by adding a $ before its name.

$ img=nc-test.fits
$ astarithmetic $img $img nan where -hINPUT-NO-SKY -hDETECTIONS      \

To invert the result (only keep the values of detected pixels), you can flip the detected pixel values (from 0 to 1 and vice-versa) by adding a not after the second $img:

$ astarithmetic $img $img not nan where -hINPUT-NO-SKY -hDETECTIONS  \

NoiseChisel can produce “Check images” to help you visualize and inspect how each step is completed. You can see all the check images it can produce with this command.

$ astnoisechisel --help | grep check

Let’s see how the quantile threshold (the first step after convolution) has been found and applied in our previous run of NoiseChisel:

$ astnoisechisel flat-ir/xdf-f160w.fits --checkqthresh

The check images/tables are also multi-extension FITS files. As you see, when check datasets are requested, NoiseChisel won’t go to the end and abort as soon as all its extensions are ready. Try listing the extensions with astfits and then opening them with ds9 as we done above. It is strongly encouraged to play with the different parameters and use the respective check images to see which step is affected by your change.

One major factor determining NoiseChisel’s the quantile threshold is NoiseChisel’s ability to identify signal in a tile (the threshold is only found on those tiles with no major signal). Therefore the larger the tiles are, number-statistics will cause less scatter, therefore helping NoiseChisel find the quantile threshold. However, if the tiles are too large, there might not be enough over the dataset or they won’t be able to deal with possible gradients. Let’s see what the default (small) tile size was with the following command.

$ astnoisechisel -P | grep tilesize

Its a 50 by 50 box. Flip back and forth between the CONVOLVED and QTHRESH_ERODE extensions of the check image to get a feeling of which tiles succeeded (have non-blank values)31. Since this is a relatively large image and we don’t have any gradients, let’s increase the tile size to 100 by 100

$ astnoisechisel flat-ir/xdf-f160w.fits --tilesize=100,100    \

Inspecting the check image, you see that there are now only a handful of useful tiles in the central parts. This shows the field is too crowded, and we should slightly decrease the tile size for a more robust result that also covers more of the dataset. Let’s set it to a 75 by 75 pixel box:

$ astnoisechisel flat-ir/xdf-f160w.fits --tilesize=75,75    \

The result seems reasonable now: we have a larger tile size than the default value, but less scatter, and the tiles cover a sufficiently wide area of the dataset. So, we’ll use this tile size for the next steps. But first, let’s clean all the temporary files and make a directory for the NoiseChisel outputs:

$ rm *.fits
$ mkdir nc
$ astnoisechisel flat-ir/xdf-f160w.fits --tilesize=75,75    \
$ astnoisechisel flat-ir/xdf-f105w.fits --tilesize=75,75    \

Before continuing with the higher-level processing of this dataset, we’ll pause to use NoiseChisel’s multi-extension output for showing how the Fits program can make it easy to work with this wonderful data container (see Fits). Let’s say you need to copy a HDU/extension from one FITS file to another. Try the command below to make an objects.fits file that contains only NoiseChisel’s binary detection map. There are similar options to conveniently cut or delete HDUs from a FITS file also.

$ astfits nc/xdf-f160w.fits --copy=DETECTIONS -odetections.fits

NoiseChisel puts some general information on its outputs in the FITS header of the respective extension. To see the full list of keywords in an extension, you can again use the Fits program like above. But instead of HDU manipulation options, give it the HDU you are interested in with -h. You can also give the HDU number (as listed in the output above), for example -h2 instead of -hDETECTIONS.

$ astfits nc/xdf-f160w.fits -hDETECTIONS

The DETSN keyword in NoiseChisel’s DETECTIONS extension contains the true pseudo-detection signal-to-noise ratio that was found by NoiseChisel on the dataset. It is not easy to find it in the middle of all the other keywords printed by the command above (especially in files that have many more keywords). To fix the problem, you can pipe the output of the command above into grep (a program for matching lines which is available on almost all Unix-like operating systems).

$ astfits nc/xdf-f160w.fits -hDETECTIONS | grep DETSN

If you just want the value of the keyword and not the full FITS keyword line, you can use AWK. In the example below, AWK will print the third word (separated by white space characters) in any line that has a first column value of DETSN.

$ astfits nc/xdf-f160w.fits -h2 | awk '$1=="DETSN" {print $3}'

The main output of NoiseChisel is the binary detection map (DETECTIONS extension), which only has two values of 1 or 0. This is useful when studying the noise, but hardly of any use when you actually want to study the targets/galaxies in the image, especially in such a deep field where the detection map of almost everything is connected. To find the galaxies over the detections, we’ll use Gnuastro’s Segment program:

$ rm *.fits
$ mkdir seg
$ astsegment nc/xdf-f160w.fits -oseg/xdf-f160w.fits

Segment’s operation is very much like NoiseChisel (in fact, prior to version 0.6, it was part of NoiseChisel), for example the output is a multi-extension FITS file, it has check images and uses the undetected regions as a reference. Please have a look at Segment’s multi-extension output with ds9 to get a good feeling of what it has done. Like NoiseChisel, the first extension is the input. The CLUMPS extension shows the true “clumps” with values that are \(\ge1\), and the diffuse regions labeled as \(-1\). In the OBJECTS extension, we see that the large detections of NoiseChisel (that may have contained many galaxies) are now broken up into separate labels. see Segment for more.

Having localized the regions of interest in the dataset, we are ready to do measurements on them with MakeCatalog. Besides the IDs, we want to measure the Right Ascension (with --ra), Declination (--dec, magnitude (--magnitude and signal-to-noise ratio (--sn) of the objects and clumps. The following command will make these measurements on Segment’s F160W output:

$ mkdir cat
$ astmkcatalog seg/xdf-f160w.fits --ids --ra --dec --magnitude --sn \
               --zeropoint=25.94                                    \
               --clumpscat --output=cat/xdf-f160w.fits

From the printed statements on the command-line, you see that MakeCatalog read all the extensions in Segment’s output for the various measurements it needed.

The output of the MakeCatalog command above is a FITS table. The two clump and object catalogs are available in the two extensions of the single FITS file32. Let’s inspect the separate extensions with the Fits program like before (as shown below). Later, we’ll inspect the table in each extension with Gnuastro’s Table program (see Table). Note that we could have used -hOBJECTS and -hCLUMPS instead of -h1 and -h2 respectively.

$ astfits  cat/xdf-f160w.fits              # Extension information
$ asttable cat/xdf-f160w.fits -h1 --info   # Objects catalog info.
$ asttable cat/xdf-f160w.fits -h1          # Objects catalog columns.
$ asttable cat/xdf-f160w.fits -h2 -i       # Clumps catalog info.
$ asttable cat/xdf-f160w.fits -h2          # Clumps catalog columns.

As you see above, when given a specific table (file name and extension), Table will print the full contents of all the columns. To see basic information about each column (for example name, units and comments), simply append a --info (or -i).

To print the contents of special column(s), just specify the column number(s) (counting from 1) or the column name(s) (if they have one). For example, if you just want the magnitude and signal-to-noise ratio of the clumps (in -h2), you can get it with any of the following commands

$ asttable cat/xdf-f160w.fits -h2 -c5,6
$ asttable cat/xdf-f160w.fits -h2 -c5,SN
$ asttable cat/xdf-f160w.fits -h2 -c5         -c6
$ asttable cat/xdf-f160w.fits -h2 -cMAGNITUDE -cSN

In the example above, the clumps catalog has two ID columns (one for the over-all clump ID and one for the ID of the clump in its host object), while the objects catalog only has one ID column. Therefore, the location of the magnitude column differs between the object and clumps catalog. So if you want to specify the columns by number, you will need to change the numbers when viewing the clump and objects catalogs. This is a useful advantage of having/using column names33.

$ asttable catalog/xdf-f160w.fits -h1 -c4 -c5
$ asttable catalog/xdf-f160w.fits -h2 -c5 -c6

Finally, the comments in MakeCatalog’s output (COMMENT keywords in the FITS headers, or lines starting with # in plain text) contain some important information about the input dataset that can be useful (for example pixel area or per-pixel surface brightness limit). For example have a look at the output of this command:

$ astfits cat/xdf-f160w.fits -h1 | grep COMMENT

To calculate colors, we also need magnitude measurements on the F105W filter. However, the galaxy properties might differ between the filters (which is the whole purpose behind measuring colors). Also, the noise properties and depth of the datasets differ. Therefore, if we simply follow the same Segment and MakeCatalog calls above for the F105W filter, we are going to get a different number of objects and clumps. Matching the two catalogs is possible (for example with Match), but the fact that the measurements will be done on different pixels, can bias the result. Since the Point spread function (PSF) of both images is very similar, an accurate color calculation can only be done when magnitudes are measured from the same pixels on both images.

The F160W image is deeper, thus providing better detection/segmentation, and redder, thus observing smaller/older stars and representing more of the mass in the galaxies. We will thus use the pixel labels generated on the F160W filter, but do the measurements on the F105W filter (using the --valuesfile option) in the command below. Notice how the only difference between this call to MakeCatalog and the previous one is --valuesfile, the value given to --zeropoint and the output name.

$ astmkcatalog seg/xdf-f160w.fits --ids --ra --dec --magnitude --sn \
               --valuesfile=nc/xdf-f105w.fits --zeropoint=26.27     \
               --clumpscat --output=cat/xdf-f105w.fits

Look into what MakeCatalog printed on the command-line. You can see that (as requested) the object and clump labels were taken from the respective extensions in seg/xdf-f160w.fits, while the values and Sky standard deviation were done on nc/xdf-f105w.fits.

Since we used the same labeled image on both filters, the number of rows in both catalogs are the same. The clumps are not affected by the hard-to-deblend and low signal-to-noise diffuse regions, they are more robust for calculating the colors (compared to objects). Therefore from this step onward, we’ll continue with clumps.

We can finally calculate the colors of the objects from these two datasets. If you inspect the contents of the two catalogs, you’ll notice that because they were both derived from the same segmentation maps, the rows are ordered identically (they correspond to the same object/clump in both filters). But to be generic (usable even when the rows aren’t ordered similarly) and display another useful program in Gnuastro, we’ll use Match.

As the name suggests, Gnuastro’s Match program will match rows based on distance (or aperture in 2D) in one (or two) columns. In the command below, the options relating to each catalog are placed under it for easy understanding. You give Match two catalogs (from the two different filters we derived above) as argument, and the HDUs containing them (if they are FITS files) with the --hdu and --hdu2 options. The --ccol1 and --ccol2 options specify which columns should be matched with which in the two catalogs. With --aperture you specify the acceptable error (radius in 2D), in the same units as the columns (see below for why we have requested an aperture of 0.35 arcseconds, or less than 6 HST pixels).

The --outcols is a very convenient feature in Match: you can use it to specify which columns from the two catalogs you want in the output. If the first character is an ‘a’, the respective matched column (number or name, similar to Table above) in the first catalog will be written in the output table. When the first character is a ‘b’, the respective column from the second catalog will be written in the output. You can use this to mix the desired matched columns from both catalogs in the output.

$ astmatch cat/xdf-f160w.fits           cat/xdf-f105w.fits         \
           --hdu=CLUMPS                 --hdu2=CLUMPS              \
           --ccol1=RA,DEC               --ccol2=RA,DEC             \
           --aperture=0.35/3600                                    \
           --outcols=a1,a2,aRA,aDEC,aMAGNITUDE,aSN,bMAGNITUDE,bSN  \
           --log --output=cat/xdf-f160w-f105w.fits

By default (when --quiet isn’t called), the Match program will just print the number of matched rows in the standard output. If you have a look at your input catalogs, this should be the same as the number of rows in them. Let’s have a look at the columns in the matched catalog:

$ asttable cat/xdf-f160w-f105w.fits -i

Indeed, its exactly the columns we wanted. There is just one confusion however: there are two MAGNITUDE and SN columns. Right now, you know that the first one was from the F160W filter, and the second was for F105W. But in one hour, you’ll start doubting your self: going through your command history, trying to answer this question: “which magnitude corresponds to which filter?”. You should never torture your future-self (or colleagues) like this! So, let’s rename these confusing columns in the matched catalog. The FITS standard for tables stores the column names in the TTYPE header keywords, so let’s have a look:

$ astfits cat/xdf-f160w-f105w.fits -h1 | grep TTYPE

Changing/updating the column names is as easy as updating the values to these options with the first command below, and with the second, confirm this change:

$ astfits cat/xdf-f160w-f105w.fits -h1                          \
          --update=TTYPE5,MAG_F160W   --update=TTYPE6,SN_F160W  \
          --update=TTYPE7,MAG_F105W   --update=TTYPE8,SN_F105W
$ asttable cat/xdf-f160w-f105w.fits -i

If you noticed, when running Match, the previous command, we also asked for --log. Many Gnuastro programs have this option to provide some detailed information on their operation in case you are curious. Here, we are using it to justify the value we gave to --aperture. Even though you asked for the output to be written in the cat directory, a listing of the contents of your current directory will show you an extra astmatch.fits file. Let’s have a look at what columns it contains.

$ ls
$ asttable astmatch.log -i

The MATCH_DIST column contains the distance of the matched rows, let’s have a look at the distribution of values in this column. You might be asking yourself “why should the positions of the two filters differ when I gave MakeCatalog the same segmentation map?” The reason is that the central positions are flux-weighted. Therefore the --valuesfile dataset you give to MakeCatalog will also affect the center measurements34. Recall that the Spectral Energy Distribution (SED) of galaxies is not flat and they have substructure, therefore, they can have different shapes/morphologies in different filters.

Gnuastro has a simple program for basic statistical analysis. The command below will print some basic information about the distribution (minimum, maximum, median and etc), along with a cute little ASCII histogram to visually help you understand the distribution on the command-line without the need for a graphic user interface (see Invoking Statistics). This ASCII histogram can be useful when you just want some coarse and general information on the input dataset. It is also useful when working on a server (where you may not have graphic user interface), and finally, its fast.

$ aststatistics astmatch.fits -cMATCH_DIST

The units of this column are the same as the columns you gave to Match: in degrees. You see that while almost all the objects matched very nicely, the maximum distance is roughly 0.31 arcseconds. This is why we asked for an aperture of 0.35 arcseconds when doing the match.

We can now use AWK to find the colors. We’ll ask AWK to only use rows that don’t have a NaN magnitude in either filter35. We will also ignore columns which don’t have reliable F105W measurement (with a S/N less than 736).

$ asttable cat/xdf-f160w-f105w.fits -cMAG_F160W,MAG_F105W,SN_F105W  \
           | awk '$1!="nan" && $2!="nan" && $3>7 {print $2-$1}'     \
           > f105w-f160w.txt

You can inspect the distribution of colors with the Statistics program again:

$ aststatistics f105w-f160w.txt -c1

You can later use Gnuastro’s Statistics program with the --histogram option to build a much more fine-grained histogram as a table to feed into your favorite plotting program for a much more accurate/appealing plot (for example with PGFPlots in LaTeX). If you just want a specific measure, for example the mean, median and standard deviation, you can ask for them specifically with this command:

$ aststatistics f105w-f160w.txt -c1 --mean --median --std

Some researchers prefer to have colors in a fixed aperture for all the objects. The colors we calculated above used a different segmentation map for each object. This might not satisfy some science cases. So, let’s make a fixed aperture catalog. To make an catalog from fixed apertures, we should make a labeled image which has a fixed label for each aperture. That labeled image can be given to MakeCatalog instead of Segment’s labeled detection image.

To generate the apertures catalog, we’ll first read the positions from F160W catalog and set the other parameters of each profile to be a fixed circle of radius 5 pixels (we want all apertures to be identical in this scenario).

$ rm *.fits *.txt
$ asttable cat/xdf-f160w.fits -hCLUMPS -cRA,DEC                    \
           | awk '!/^#/{print NR, $1, $2, 5, 5, 0, 0, 1, NR, 1}' \
           > apertures.txt

We can now feed this catalog into MakeProfiles to build the apertures for us. See Invoking MakeProfiles for a description of the options. The most important for this particular job is --mforflatpix, it tells MakeProfiles that the values in the magnitude column should be used for each pixel of a flat profile. Without it, MakeProfiles would build the profiles such that the sum of the pixels of each profile would have a magnitude (in log-scale) of the value given in that column (what you would expect when simulating a galaxy for example).

$ astmkprof apertures.txt --background=flat-ir/xdf-f160w.fits     \
            --clearcanvas --replace --type=int16 --mforflatpix    \

The first thing you might notice in the printed information is that the profiles are not built in order. This is because MakeProfiles works in parallel, and parallel CPU operations are asynchronous. You can try running MakeProfiles with one thread (using --numthreads=1) to see how order is respected in that case.

Open the output apertures.fits file and see the result. Where the apertures overlap, you will notice that one label has replaced the other (because of the --replace option). In the future, MakeCatalog will be able to work with overlapping labels, but currently it doesn’t. If you are interested, please join us in completing Gnuastro with added improvements like this (see task 14750 37).

We can now feed the apertures.fits labeled image into MakeCatalog instead of Segment’s output as shown below. In comparison with the previous MakeCatalog call, you will notice that there is no more --clumpscat option, since each aperture is treated as a separate “object” here.

$ astmkcatalog apertures.fits -h1 --zeropoint=26.27        \
               --valuesfile=nc/xdf-f105w.fits              \
               --ids --ra --dec --magnitude --sn           \

This catalog has the same number of rows as the catalog produced from clumps, therefore similar to how we found colors, you can compare the aperture and clump magnitudes for example. You can also change the filter name and zeropoint magnitudes and run this command again to have the fixed aperture magnitude in the F160W filter and measure colors on apertures.

As a final step, let’s go back to the original clumps-based catalogs we generated before. We’ll find the objects with the strongest color and make a cutout to inspect them visually and finally, we’ll see how they are located on the image.

First, let’s see what the objects with a color more than two magnitudes look like. As you see, this is very much like the command above for selecting the colors, only instead of printing the color, we’ll print the RA and Dec. With the command below, the positions of all lines with a color more than 1.5 will be put in reddest.txt

$ asttable cat/xdf-f160w-f105w.fits                                \
           -cMAG_F160W,MAG_F105W,SN_F105W,RA,DEC                   \
           | awk '$1!="nan" && $2!="nan" && $2-$1>1.5 && $3>7      \
                  {print $4,$5}' > reddest.txt

We can now feed reddest.txt into Gnuastro’s crop to see what these objects look like. To keep things clean, we’ll make a directory called crop-red and ask Crop to save the crops in this directory. We’ll also add a -f160w.fits suffix to the crops (to remind us which image they came from). The width of the crops will be 15 arcseconds.

$ mkdir crop-red
$ astcrop --mode=wcs --coordcol=3 --coordcol=4 flat-ir/xdf-f160w.fits \
          --catalog=reddest.txt --width=15/3600,15/3600               \
          --suffix=-f160w.fits --output=crop-red

Like the MakeProfiles command above, you might notice that the crops aren’t made in order. This is because each crop is independent of the rest, therefore crops are done in parallel, and parallel operations are asynchronous. In the command above, you can change f160w to f105w to make the crops in both filters.

To view the crops more easily (not having to open ds9 for each image), you can convert the FITS crops into the JPEG format with a shell loop like below.

$ cd crop-red
$ for f in *.fits; do                                                  \
    astconvertt $f --fluxlow=-0.001 --fluxhigh=0.005 --invert -ojpg;   \
$ cd ..

You can now use your general graphic user interface image viewer to flip through the images more easily. On GNOME, you can use the “Eye of GNOME” image viewer (with executable name of eog). Run the command below to open the first one (if you aren’t using GNOME, use the command of your image viewer instead of eog):

$ eog 1-f160w.jpg

In Eye of GNOME, you can flip through the images and compare them visually more easily by pressing the <SPACE> key. Of course, the flux ranges have been chosen generically here for seeing the fainter parts. Therefore, brighter objects will be fully black.

The for loop above to convert the images will do the job in series: each file is converted only after the previous one is complete. If you have GNU Parallel, you can greatly speed up this conversion. GNU Parallel will run the separate commands simultaneously on different CPU threads in parallel. For more information on efficiently using your threads, see Multi-threaded operations. Here is a replacement for the shell for loop above using GNU Parallel.

$ cd crop-red
$ parallel astconvertt --fluxlow=-0.001 --fluxhigh=0.005 --invert   \
           -ojpg ::: *.fits
$ cd ..

As the final action, let’s see how these objects are positioned over the dataset. DS9 has the “Region”s concept for this purpose. You just have to convert your catalog into a “region file” to feed into DS9. To do that, you can use AWK again as shown below.

$ awk 'BEGIN{print "# Region file format: DS9 version 4.1";     \
             print "global color=green width=2";                \
             print "fk5";}                                      \
       {printf "circle(%s,%s,1\")\n", $1, $2;}' reddest.txt     \
       > reddest.reg

This region file can be loaded into DS9 with its -regions option to display over any image (that has world coordinate system). In the example below, we’ll open Segment’s output and load the regions over all the extensions (to see the image and the respective clump):

$ ds9 -mecube seg/xdf-f160w.fits -zscale -zoom to fit    \
      -regions load all reddest.reg

In conclusion, we hope this extended tutorial has been a good starting point to help in your exciting research. If this book or any of the programs in Gnuastro have been useful for your research, please cite the respective papers and share your thoughts and suggestions with us (it can be very encouraging). All Gnuastro programs have a --cite option to help you cite the authors’ work more easily. Just note that it may be necessary to cite additional papers for different programs, so please try it out for any program you used.

$ astmkcatalog --cite
$ astnoisechisel --cite

Next: , Previous: , Up: Tutorials   [Contents][Index]

2.3 Detecting large extended targets

The outer wings of large and extended objects can sink into the noise very gradually and can have a large variety of shapes (for example due to tidal interactions). Therefore separating the outer boundaries of the galaxies from the noise can be particularly tricky. Besides causing an under-estimation in the total estimated brightness of the target, failure to detect such faint wings will also cause a bias in the noise measurements, thereby hampering the accuracy of any measurement on the dataset. Therefore even if they don’t constitute a significant fraction of the target’s light, or aren’t your primary target, these regions must not be ignored. In this tutorial, we’ll walk you through the strategy of detecting such targets using NoiseChisel.

Don’t start with this tutorial: If you haven’t already completed General program usage tutorial, we strongly recommend going through that tutorial before starting this one. Basic features like access to this book on the command-line, the configuration files of Gnuastro’s programs, benefiting from the modular nature of the programs, viewing multi-extension FITS files, or using NoiseChisel’s outputs are discussed in more detail there.

We’ll try to detect the faint tidal wings of the beautiful M51 group38 in this tutorial. We’ll use a dataset/image from the public Sloan Digital Sky Survey, or SDSS. Due to its more peculiar low surface brightness structure/features, we’ll focus on the dwarf companion galaxy of the group (or NGC 5195). To get the image, you can use SDSS’s Simple field search tool. As long as it is covered by the SDSS, you can find an image containing your desired target either by providing a standard name (if it has one), or its coordinates. To access the dataset we will use here, write NGC5195 in the “Object Name” field and press “Submit” button.

Type the example commands: Try to type the example commands on your terminal and use the history feature of your command-line (by pressing the “up” button to retrieve previous commands). Don’t simply copy and paste the commands shown here. This will help simulate future situations when you are processing your own datasets.

You can see the list of available filters under the color image. For this demonstration, we’ll use the r-band filter image. By clicking on the “r-band FITS” link, you can download the image. Alternatively, you can just run the following command to download it with GNU Wget39. To keep things clean, let’s also put it in a directory called ngc5195. With the -O option, we are asking Wget to save the downloaded file with a more manageable name: r.fits.bz2 (this is an r-band image of NGC 5195, which was the directory name).

$ mkdir ngc5195
$ cd ngc5195
$ topurl=
$ wget $topurl/301/3716/6/frame-r-003716-6-0117.fits.bz2 -Or.fits.bz2

This server keeps the files in a Bzip2 compressed file format. So we’ll first decompress it with the following command. By convention, compression programs delete the original file (compressed when uncompressing, or uncompressed when compressing). To keep the original file, you can use the --keep or -k option which is available in most compression programs for this job. Here, we don’t need the compressed file any more, so we’ll just let bunzip delete it for us and keep the directory clean.

$ bunzip2 r.fits.bz2

Let’s see how NoiseChisel operates on it with its default parameters:

$ astnoisechisel r.fits -h0

As described in NoiseChisel output, NoiseChisel’s default output is a multi-extension FITS file. A method to view them effectively and easily is discussed in Viewing multiextension FITS images.

Open the output r_detected.fits file and you will immediately notice how NoiseChisel’s default configuration is not suitable for this dataset: the Sky estimation has failed so terribly that the tile grid (where the Sky was estimated, and subtracted) is visible in the first extension (input dataset subtracted by the Sky value). If you look into the third and fourth extensions (the Sky and its standard deviation) you will see how they exactly map NGC 5195! This is not good! There shouldn’t be any signature of your extended target on the Sky and its standard deviation images. After all, the Sky is suppose to be the average value in the absence of signal, see Sky value.

The fact that signal has been detected as Sky shows that you haven’t done a good detection. Generally, any time your target is much larger than the tile size and the signal is almost flat (like this case), this will happen, even if it isn’t dramatic enough to be seen in the first extension. Therefore, the best place to check the accuracy of your detection is the noise extensions (third and fourth extension) of NoiseChisel’s output.

When dominated by the background, noise has a symmetric distribution. However, signal is not symmetric (we don’t have negative signal). Therefore when non-constant signal is present in a noisy dataset, the distribution will be positively skewed. This skewness is a good measure of how much signal we have in the distribution. The skewness can be accurately measured by the difference in the mode and median, for more see Quantifying signal in a tile, and Appendix C Akhlaghi and Ichikawa [2015].

Skewness is only a proxy for signal when the signal has structure (varies per pixel). Therefore, when it is approximately constant over a whole tile, or sub-set of the image, the signal’s effect is just to shift the symmetric center of the noise distribution to the positive and there won’t be any skewness: this positive40 shift that preserves the symmetric distribution is the Sky value. When there is a gradient over the dataset, different tiles will have different constant shifts/Sky-values, for example see Figure 11 of Akhlaghi and Ichikawa [2015].

To get less scatter in measuring the mode and median (and thus better estimate the skewness), you will need a larger tile. In Gnuastro, you can see the option values (--tilesize in this case) by adding the -P option to your last command. Try it. You can clearly see that the default tile size is indeed much smaller than this (huge) dwarf galaxy. Therefore NoiseChisel was unable to identify the skewness within the tiles under NGC 5159. Recall that NoiseChisel only uses tiles with no signal/skewness to define its threshold. Because of this, the threshold has been over-estimated on those tiles and further exacerbated the non-detection of the diffuse regions. To see which tiles were used for estimating the quantile threshold (no skewness was measured), you can use NoiseChisel’s --checkqthresh option:

$ astnoisechisel r.fits -h0 --checkqthresh

Notice how this option doesn’t allow NoiseChisel to finish. NoiseChisel aborted after finding the quantile thresholds. When you call any of NoiseChisel’s --check* options, by default, it will abort as soon as all the check steps have been written in the check file (a multi-extension FITS file). To optimize the threshold-related settings for this image, we’ll be playing with this tile for the majority of this tutorial. So let’s have a closer look at it.

The first extension of r_qthresh.fits (CONVOLVED) is the convolved input image (where the threshold is defined and applied), for more on the effect of convolution and thresholding, see Sections 3.1.1 and 3.1.2 of Akhlaghi and Ichikawa [2015]. The second extension (QTHRESH_ERODE) has a blank value for all the pixels of any tile that was identified as having significant signal. Playing a little with the dynamic range of this extension, you clearly see how the non-blank tiles around NGC 5195 have a gradient. You do not want this behavior. The ultimate purpose of the next few trials will be to remove the gradient from the non-blank tiles.

The next two extensions (QTHRESH_NOERODE and QTHRESH_EXPAND) are for later steps in NoiseChisel. The same tiles are masked, but those with a value, have a different value compared to QTHRESH_ERODE. In the subsequent three extensions, you can see how the blank tiles are filled/interpolated. The subsequent three extensions show the smoothed tile values. Finally in the last extension (QTHRESH-APPLIED), you can see the effect of applying QTHRESH_ERODE on CONVOLVED (pixels with a value of 0 were below the threshold).

Skipping convolution for faster tests: The slowest step of NoiseChisel is the convolution of the input dataset. Therefore when your dataset is large (unlike the one in this test), and you are not changing the input dataset or kernel in multiple runs (as in the tests of this tutorial), it is faster to do the convolution separately once (using Convolve) and use NoiseChisel’s --convolved option to directly feed the convolved image and avoid convolution. For more on --convolved, see NoiseChisel input.

Fortunately this image is large and has a nice and clean region also (filled with very small/distant stars and galaxies). So our first solution is to increase the tile size. To identify the skewness caused by NGC 5195 on the tiles under it, we thus have to choose a tile size that is larger than the gradient of the signal. Let’s try a 100 by 100 tile size:

$ astnoisechisel r.fits -h0 --tilesize=100,100 --checkqthresh

You can clearly see the effect of this increased tile size: they are much larger. Nevertheless the higher valued tiles are still systematically surrounding NGC 5195. As a result, when flipping through the interpolated and smoothed values, you can still see the signature of the galaxy, and the ugly tile signatures are still present in QTHRESH-APPLIED. So let’s increase the tile size even further (check the result of the first before going to the second):

$ astnoisechisel r.fits -h0 --tilesize=150,150 --checkqthresh
$ astnoisechisel r.fits -h0 --tilesize=200,200 --checkqthresh

The number of tiles with a gradient does indeed decrease with a larger tile size, but you still see a gradient in the raw values. The tile signatures in the thresholded image are also still present. These are not a good sign. So, let’s keep the 200 by 200 tile size and start playing with the other constraint that we have: the acceptable distance (in quantile), between the mode and median.

The tile size is now very large (16 times the area of the default tile size). We thus have much less scatter in the estimation of the mode and median and we can afford to decrease the acceptable distance between the two. The acceptable distance can be set through the --modmedqdiff option (read as “mode-median quantile difference”). Before trying the next command, run the previous command with a -P to see what value it originally had.

$ astnoisechisel r.fits -h0 --tilesize=200,200 --modmedqdiff=0.005    \

But this command doesn’t finish like the previous ones. A r_qthresh.fits file was created, but instead of saying that the quantile thresholds have been applied, a long error message is printed. Please read the error message before continuing to read here.

The error message fully describes the problem and even proposes solutions. As suggested there, the ideal solution would be to use SDSS images outside of this field and add them to this one to have a larger input image (with a larger area outside the diffuse regions). But we don’t always have this luxury, so let’s keep using to this image for the rest of this tutorial.

First, open r_qthresh.fits and have a look at the successful tiles. Unlike the previous --checkqthresh outputs, this one only has four extensions: as the error message explains, the interpolation (to give values to blank tiles) has not been done. Therefore its check results aren’t present.

At the start of the error message, NoiseChisel tells us how many tiles passed the test for having no significant signal: six. Looking closely at the dataset, we see that outside NGC 5195, there is no background gradient (the background is a fixed value). Our tile sizes are also very large (and thus don’t have much scatter). So a good way to loosen up the parameters can be to simply decrease the number of neighboring tiles needed for interpolation, with --interpnumngb (read as “interpolation number of neighbors”).

$ astnoisechisel r.fits -h0 --tilesize=200,200 --modmedqdiff=0.005    \
                 --interpnumngb=6 --checkqthresh

There is no longer any significant gradient in the usable tiles and no signature of NGC 5195 exists in the interpolated and smoothed values. But before finishing the quantile threshold, let’s have a closer look at the thresholded image in QTHRESH-APPLIED. Slide the dynamic range in your FITS viewer so 0 valued pixels are black and all non-zero pixels are white. You will see that the black holes are not evenly distributed. Those that follow the tail of NGC 5195 are systematically smaller than those in the far-right of the image. This suggests that we can decrease the quantile threshold (--qthresh) even further: there is still signal down there!

$ rm r_qthresh.fits
$ astnoisechisel r.fits -h0 --tilesize=200,200 --modmedqdiff=0.005    \
                 --interpnumngb=6 --qthresh=0.2

Since the quantile threshold of the previous command was satisfactory, we finally removed --checkqthresh to let NoiseChisel proceed until completion. Looking at the DETECTIONS extension of NoiseChisel’s output, we see the right-ward edges in particular have many holes that are fully surrounded by signal and the signal stretches out in the noise very thinly. This suggests that there is still signal that can be detected. So we’ll decrease the growth quantile (for larger/deeper growth into the noise, with --detgrowquant) and increase the size of holes that can be filled (if they are fully surrounded by signal, with --detgrowmaxholesize). Since we are done with our detection, to facilitate later steps, we’ll also add the --label option so the connected regions get different labels.

$ astnoisechisel r.fits -h0 --tilesize=200,200 --modmedqdiff=0.005    \
                 --interpnumngb=6 --qthresh=0.2 --detgrowquant=0.6    \
                 --detgrowmaxholesize=10000 --label

Looking into the output, we now clearly see that the tidal features of M51 and NGC 5195 are detected nicely in the same direction as expected (towards the bottom right side of the image). However, as discussed above, the best measure of good detection is the noise, not the detections themselves. So let’s look at the Sky and its Standard deviation. The Sky standard deviation no longer has any footprint of NGC 5195. But the Sky still has a very faint shadow of the dwarf galaxy (the values on the left are larger than on the right). However, this gradient in the Sky (output of first command below) is much less (by \(\sim20\) times) than the standard deviation (output of second command). So we can stop playing with NoiseChisel here, and leave the configuration for a more accurate detection to you.

$ aststatistics r_detected.fits -hSKY --maximum --minimum       \
                | awk '{print $1-$2}'
$ aststatistics r_detected.fits -hSKY_STD --mean

Let’s see how deeply/successfully we carved out M51 and NGC 5195’s tail from the noise. For this measurement, we’ll need to estimate the average flux on the outer edges of the detection. Fortunately all this can be done with a few simple commands (and no higher-level language mini-environments) using Arithmetic and MakeCatalog.

The M51 group detection is by far the largest detection in this image. We can thus easily find the ID/label that corresponds to it. We’ll first run MakeCatalog to find the area of all the detections, then we’ll use AWK to find the ID of the largest object and keep it as a shell variable (id):

$ astmkcatalog r_detected.fits --ids --geoarea -hDETECTIONS -ocat.txt
$ id=$(awk '!/^#/{if($2>max) {id=$1; max=$2}} END{print id}' cat.txt)

To separate the outer edges of the detections, we’ll need to “erode” the detections. We’ll erode two times (one time would be too thin for such a huge object), using a maximum connectivity of 2 (8-connected neighbors). We’ll then save the output in eroded.fits.

$ astarithmetic r_detected.fits 0 gt 2 erode -hDETECTIONS -oeroded.fits

We should now just keep the pixels that have the ID of the M51 group, but a value of 0 in erode.fits. We’ll keep the output in boundary.fits.

$ astarithmetic r_detected.fits $id eq eroded.fits 0 eq and     \
                -hDETECTIONS -h1 -oboundary.fits

Open the image and have a look. You’ll see that the detected edge of the M51 group is now clearly visible. You can use boundary.fits to mark (set to blank) this boundary on the input image and get a visual feeling of how far it extends:

$ astarithmetic r.fits boundary.fits nan where -ob-masked.fits -h0

To quantify how deep we have detected the low-surface brightness regions, we’ll use the command below. In short it just divides all the 1-valued pixels of boundary.fits in the Sky subtracted input (first extension of NoiseChisel’s output) by the pixel standard deviation of the same pixel. This will give us a signal-to-noise ratio image. The mean value of this image shows the level of surface brightness that we have achieved.

You can also break the command below into multiple calls to Arithmetic and create temporary files to understand it better. However, if you have a look at Reverse polish notation and Arithmetic operators, you should be able to easily understand what your computer does when you run this command41.

$ astarithmetic r_detected.fits boundary.fits not nan where \
                r_detected.fits /                           \
                meanvalue                                   \
                -hINPUT-NO-SKY -h1 -hSKY_STD --quiet
--> 0.0511864

The outer wings where therefore non-parametrically detected until \(\rm{S/N}\approx0.05\).

In interpreting this value, you should just have in mind that NoiseChisel works based on the contiguity of signal in the pixels. Therefore the larger the object, the deeper NoiseChisel can carve it out of the noise. In other words, this reported depth, is only for this particular object and dataset, processed with this particular NoiseChisel configuration: if the M51 group in this image was larger/smaller than this, or if the image was larger/smaller, or if we had used a different configuration, we would go deeper/shallower.

The NoiseChisel configuration found here is NOT GENERIC for any large object: As you saw above, the reason we chose this particular configuration for NoiseChisel to detect the wings of the M51 group was strongly influenced by this particular object in this particular image. When low surface brightness signal takes over such a large fraction of your dataset (and you want to accurately detect/account for it), to make sure that it is successfully detected, you will need some manual checking, intervention, or customization. In other words, to make sure that your noise measurements are least affected by the signal42.

To avoid typing all these options every time you run NoiseChisel on this image, you can use Gnuastro’s configuration files, see Configuration files. For an applied example of setting/using them, see General program usage tutorial.

To continue your analysis of such datasets with extended emission, you can use Segment to identify all the “clumps” over the diffuse regions: background galaxies and foreground stars.

$ astsegment r_detected.fits

Open the output r_detected_segmented.fits as a multi-extension data cube like before and flip through the first and second extensions to see the detected clumps (all pixels with a value larger than 1). To optimize the parameters and make sure you have detected what you wanted, its highly recommended to visually inspect the detected clumps on the input image.

For visual inspection, you can make a simple shell script like below. It will first call MakeCatalog to estimate the positions of the clumps, then make an SAO ds9 region file and open ds9 with the image and region file. Recall that in a shell script, the numeric variables (like $1, $2, and $3 in the example below) represent the arguments given to the script. But when used in the AWK arguments, they refer to column numbers.

To create the shell script, using your favorite text editor, put the contents below into a file called Recall that everything after a # is just comments to help you understand the command (so read them!). Also note that if you are copying from the PDF version of this book, fix the single quotes in the AWK command.

#! /bin/bash
set -e	   # Stop execution when there is an error.
set -u	   # Stop execution when a variable is not initialized.

# Run MakeCatalog to write the coordinates into a FITS table.
# Default output is `$1_cat.fits'.
astmkcatalog $1.fits --clumpscat --ids --ra --dec

# Use Gnuastro's Table program to read the RA and Dec columns of the
# clumps catalog (in the `CLUMPS' extension). Then pipe the columns
# to AWK for saving as a DS9 region file.
asttable $1"_cat.fits" -hCLUMPS -cRA,DEC                               \
         | awk 'BEGIN { print "# Region file format: DS9 version 4.1"; \
                        print "global color=green width=1";            \
                        print "fk5" }                                  \
                { printf "circle(%s,%s,1\")\n", $1, $2 }' > $1.reg

# Show the image (with the requested color scale) and the region file.
ds9 -geometry 1800x3000 -mecube $1.fits -zoom to fit                   \
    -scale limits $2 $3 -regions load all $1.reg

# Clean up (delete intermediate files).
rm $1"_cat.fits" $1.reg

Finally, you just have to activate the script’s executable flag with the command below. This will enable you to directly/easily call the script as a command.

$ chmod +x

This script doesn’t expect the .fits suffix of the input’s filename as the first argument. Because the script produces intermediate files (a catalog and DS9 region file, which are later deleted). However, we don’t want multiple instances of the script (on different files in the same directory) to collide (read/write to the same intermediate files). Therefore, we have used suffixes added to the input’s name to identify the intermediate files. Note how all the $1 instances in the commands (not within the AWK command43) are followed by a suffix. If you want to keep the intermediate files, put a # at the start of the last line.

The few, but high-valued, bright pixels in the central parts of the galaxies can hinder easy visual inspection of the fainter parts of the image. With the second and third arguments to this script, you can set the numerical values of the color map (first is minimum/black, second is maximum/white). You can call this script with any44 output of Segment (when --rawoutput is not used) with a command like this:

$ ./ r_detected_segmented -0.1 2

Go ahead and run this command. You will see the intermediate processing being done and finally it opens SAO DS9 for you with the regions superimposed on all the extensions of Segment’s output. The script will only finish (and give you control of the command-line) when you close DS9. If you need your access to the command-line before closing DS9, add a & after the end of the command above.

While DS9 is open, slide the dynamic range (values for black and white, or minimum/maximum values in different color schemes) and zoom into various regions of the M51 group to see if you are satisfied with the detected clumps. Don’t forget that through the “Cube” window that is opened along with DS9, you can flip through the extensions and see the actual clumps also. The questions you should be asking your self are these: 1) Which real clumps (as you visually feel) have been missed? In other words, is the completeness good? 2) Are there any clumps which you feel are false? In other words, is the purity good?

Note that completeness and purity are not independent of each other, they are anti-correlated: the higher your purity, the lower your completeness and vice-versa. You can see this by playing with the purity level using the --snquant option. Run Segment as shown above again with -P and see its default value. Then increase/decrease it for higher/lower purity and check the result as before. You will see that if you want the best purity, you have to sacrifice completeness and vice versa.

One interesting region to inspect in this image is the many bright peaks around the central parts of M51. Zoom into that region and inspect how many of them have actually been detected as true clumps. Do you have a good balance between completeness and purity? Also look out far into the wings of the group and inspect the completeness and purity there.

An easier way to inspect completeness (and only completeness) is to mask all the pixels detected as clumps and visually inspecting the rest of the pixels. You can do this using Arithmetic in a command like below. For easy reading of the command, we’ll define the shell variable i for the image name and save the output in masked.fits.

$ i=r_detected_segmented.fits
$ astarithmetic $i $i 0 gt nan where -hINPUT -hCLUMPS -omasked.fits

Inspecting masked.fits, you can see some very diffuse peaks that have been missed, especially as you go farther away from the group center and into the diffuse wings. This is due to the fact that with this configuration, we have focused more on the sharper clumps. To put the focus more on diffuse clumps, you can use a wider convolution kernel. Using a larger kernel can also help in detecting the existing clumps to fainter levels (thus better separating them from the surrounding diffuse signal).

You can make any kernel easily using the --kernel option in MakeProfiles. But note that a larger kernel is also going to wash-out many of the sharp/small clumps close to the center of M51 and also some smaller peaks on the wings. Please continue playing with Segment’s configuration to obtain a more complete result (while keeping reasonable purity). We’ll finish the discussion on finding true clumps at this point.

The properties of the background objects can then easily be measured using MakeCatalog. To measure the properties of the background objects (detected as clumps over the diffuse region), you shouldn’t mask the diffuse region. When measuring clump properties with MakeCatalog, the ambient flux (from the diffuse region) is calculated and subtracted. If the diffuse region is masked, its effect on the clump brightness cannot be calculated and subtracted.

To keep this tutorial short, we’ll stop here. See General program usage tutorial and Segment for more on using Segment, producing catalogs with MakeCatalog and using those catalogs.

Finally, if this book or any of the programs in Gnuastro have been useful for your research, please cite the respective papers and share your thoughts and suggestions with us (it can be very encouraging). All Gnuastro programs have a --cite option to help you cite the authors’ work more easily. Just note that it may be necessary to cite additional papers for different programs, so please use --cite with any program that has been useful in your work.

$ astmkcatalog --cite
$ astnoisechisel --cite

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2.4 Hubble visually checks and classifies his catalog

In 1924 Hubble45 announced his discovery that some of the known nebulous objects are too distant to be within the the Milky Way (or Galaxy) and that they were probably distant Galaxies46 in their own right. He had also used them to show that the redshift of the nebulae increases with their distance. So now he wants to study them more accurately to see what they actually are. Since they are nebulous or amorphous, they can’t be modeled (like stars that are always a point) easily. So there is no better way to distinguish them than to visually inspect them and see if it is possible to classify these nebulae or not.

Hubble has stored all the FITS images of the objects he wants to visually inspect in his /mnt/data/images directory. He has also stored his catalog of extra-galactic nebulae in /mnt/data/catalogs/extragalactic.txt. Any normal user on his GNU/Linux system (including himself) only has read access to the contents of the /mnt/data directory. He has done this by running this command as root:

# chmod -R 755 /mnt/data

Hubble has done this intentionally to avoid mistakenly deleting or modifying the valuable images he has taken at Mount Wilson while he is working as an ordinary user. Retaking all those images and data is simply not an option. In fact they are also in another hard disk (/dev/sdb1). So if the hard disk which stores his GNU/Linux distribution suddenly malfunctions due to work load, his data is not in harms way. That hard disk is only mounted to this directory when he wants to use it with the command:

# mount /dev/sdb1 /mnt/data

In short, Hubble wants to keep his data safe and fortunately by default Gnuastro allows for this. Hubble creates a temporary visualcheck directory in his home directory for this check. He runs the following commands to make the directory and change to it47:

$ mkdir ~/visualcheck
$ cd ~/visualcheck
$ pwd
$ ls

Hubble has multiple images in /mnt/data/images, some of his targets might be on the edges of an image and so several images need to be stitched to give a good view of them. Also his extra-galactic targets belong to various pointings in the sky, so they are not in one large image. Gnuastro’s Crop is just the program he wants. The catalog in extragalactic.txt is a plain text file which stores the basic information of all his known 200 extra-galactic nebulae. If you don’t have any particular catalog and accompanying image, you can use one the Hubble Space Telescope F160W catalog that we produced in General program usage tutorial along with the accompanying image (specify the exact image name, not /mnt/data/images/*.fits). You can select the brightest galaxies for an easier classification.

In its second column, the catalog has each object’s Right Ascension (the first column is a label he has given to each object), and in the third, the object’s declination (which he specifies with the --coordcol option). Also, since the coordinates are in the world coordinate system (WCS, not pixel positions) units, he adds --mode=wcs.

$ astcrop --coordcol=2 --coordcol=3 /mnt/data/images/*.fits     \
          --mode=wcs /mnt/data/catalogs/extragalactic.txt
Crop started on Tue Jun  14 10:18:11 1932
  ---- ./4_crop.fits                  1 1
  ---- ./2_crop.fits                  1 1
  ---- ./1_crop.fits                  1 1
[[[ Truncated middle of list ]]]
  ---- ./198_crop.fits                1 1
  ---- ./195_crop.fits                1 1
  - 200 images created.
  - 200 were filled in the center.
  - 0 used more than one input.
Crop finished in:  2.429401 (seconds)

Hubble already knows that thread allocation to the the CPU cores is asynchronous. Hence each time you run it, the order of which job gets done first differs. When using Crop the order of outputs is irrelevant since each crop is independent of the rest. This is why the crops are not necessarily created in the same input order. He is satisfied with the default width of the outputs (which he inspected by running $ astcrop -P). If he wanted a different width for the cropped images, he could do that with the --wwidth option which accepts a value in arc-seconds. When he lists the contents of the directory again he finds his 200 objects as separate FITS images.

$ ls
1_crop.fits 2_crop.fits ... 200_crop.fits

The FITS image format was not designed for efficient/fast viewing, but mainly for accurate storing of the data. So he chooses to convert the cropped images to a more common image format to view them more quickly and easily through standard image viewers (which load much faster than FITS image viewer). JPEG is one of the most recognized image formats that is supported by most image viewers. Fortunately Gnuastro has just such a tool to convert various types of file types to and from each other: ConvertType. Hubble has already heard of GNU Parallel from one of his colleagues at Mount Wilson Observatory. It allows multiple instances of a command to be run simultaneously on the system, so he uses it in conjunction with ConvertType to convert all the images to JPEG.

$ parallel astconvertt -ojpg ::: *_crop.fits

For his graphical user interface Hubble is using GNOME which is the default in most distributions in GNU/Linux. The basic image viewer in GNOME is the Eye of GNOME, which has the executable file name eog 48. Since he has used it before, he knows that once it opens an image, he can use the ENTER or SPACE keys on the keyboard to go to the next image in the directory or the Backspace key to go the previous image. So he opens the image of the first object with the command below and with his cup of coffee in his other hand, he flips through his targets very fast to get a good initial impression of the morphologies of these extra-galactic nebulae.

$ eog 1_crop.jpg

Hubble’s cup of coffee is now finished and he also got a nice general impression of the shapes of the nebulae. He tentatively/mentally classified the objects into three classes while doing the visual inspection. One group of the nebulae have a very simple elliptical shape and seem to have no internal special structure, so he gives them code 1. Another clearly different class are those which have spiral arms which he associates with code 2 and finally there seems to be a class of nebulae in between which appear to have a disk but no spiral arms, he gives them code 3.

Now he wants to know how many of the nebulae in his extra-galactic sample are within each class. Repeating the same process above and writing the results on paper is very time consuming and prone to errors. Fortunately Hubble knows the basics of GNU Bash shell programming, so he writes the following short script with a loop to help him with the job. After all, computers are made for us to operate and knowing basic shell programming gives Hubble this ability to creatively operate the computer as he wants. So using GNU Emacs49 (his favorite text editor) he puts the following text in a file named

for name in *.jpg
    eog $name &
    echo -n "$name belongs to class: "
    read class
    echo $name $class >> classified.txt
    kill $processid

Fortunately GNU Emacs or even simpler editors like Gedit (part of the GNOME graphical user interface) will display the variables and shell constructs in different colors which can really help in understanding the script. Put simply, the for loop gets the name of each JPEG file in the directory this script is run in and puts it in name. In the shell, the value of a variable is used by putting a $ sign before the variable name. Then Eye of GNOME is run on the image in the background to show him that image and its process ID is saved internally (this is necessary to close Eye of GNOME later). The shell then prompts the user to specify a class and after saving it in class, it prints the file name and the given class in the next line of a file named classified.txt. To make the script executable (so he can run it later any time he wants) he runs:

$ chmod +x

Now he is ready to do the classification, so he runs the script:

$ ./

In the end he can delete all the JPEG and FITS files along with Crop’s log file with the following short command. The only files remaining are the script and the result of the classification.

$ rm *.jpg *.fits astcrop.txt
$ ls

He can now use classified.txt as input to a plotting program to plot the histogram of the classes and start making interpretations about what these nebulous objects that are outside of the Galaxy are.

Next: , Previous: , Up: Top   [Contents][Index]

3 Installation

The latest released version of Gnuastro source code is always available at the following URL:

Quick start describes the commands necessary to configure, build, and install Gnuastro on your system. This chapter will be useful in cases where the simple procedure above is not sufficient, for example your system lacks a mandatory/optional dependency (in other words, you can’t pass the $ ./configure step), or you want greater customization, or you want to build and install Gnuastro from other random points in its history, or you want a higher level of control on the installation. Thus if you were happy with downloading the tarball and following Quick start, then you can safely ignore this chapter and come back to it in the future if you need more customization.

Dependencies describes the mandatory, optional and bootstrapping dependencies of Gnuastro. Only the first group are required/mandatory when you are building Gnuastro using a tarball (see Release tarball), they are very basic and low-level tools used in most astronomical software, so you might already have them installed, if not they are very easy to install as described for each. Downloading the source discusses the two methods you can obtain the source code: as a tarball (a significant snapshot in Gnuastro’s history), or the full history50. The latter allows you to build Gnuastro at any random point in its history (for example to get bug fixes or new features that are not released as a tarball yet).

The building and installation of Gnuastro is heavily customizable, to learn more about them, see Build and install. This section is essentially a thorough explanation of the steps in Quick start. It discusses ways you can influence the building and installation. If you encounter any problems in the installation process, it is probably already explained in Known issues. In Other useful software the installation and usage of some other free software that are not directly required by Gnuastro but might be useful in conjunction with it is discussed.

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3.1 Dependencies

A minimal set of dependencies are mandatory for building Gnuastro from the standard tarball release. If they are not present you cannot pass Gnuastro’s configuration step. The mandatory dependencies are therefore very basic (low-level) tools which are easy to obtain, build and install, see Mandatory dependencies for a full discussion.

If you have the packages of Optional dependencies, Gnuastro will have additional functionality (for example converting FITS images to JPEG or PDF). If you are installing from a tarball as explained in Quick start, you can stop reading after this section. If you are cloning the version controlled source (see Version controlled source), an additional bootstrapping step is required before configuration and its dependencies are explained in Bootstrapping dependencies.

Your operating system’s package manager is an easy and convenient way to download and install the dependencies that are already pre-built for your operating system. In Dependencies from package managers, we’ll list some common operating system package manager commands to install the optional and mandatory dependencies.

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3.1.1 Mandatory dependencies

The mandatory Gnuastro dependencies are very basic and low-level tools. They all follow the same basic GNU based build system (like that shown in Quick start), so even if you don’t have them, installing them should be pretty straightforward. In this section we explain each program and any specific note that might be necessary in the installation.

Next: , Previous: , Up: Mandatory dependencies   [Contents][Index] GNU Scientific library

The GNU Scientific Library, or GSL, is a large collection of functions that are very useful in scientific applications, for example integration, random number generation, and Fast Fourier Transform among many others. To install GSL from source, you can run the following commands after you have downloaded gsl-latest.tar.gz:

$ tar xf gsl-latest.tar.gz
$ cd gsl-X.X                     # Replace X.X with version number.
$ ./configure
$ make -j8                       # Replace 8 with no. CPU threads.
$ make check
$ sudo make install

Next: , Previous: , Up: Mandatory dependencies   [Contents][Index] CFITSIO

CFITSIO is the closest you can get to the pixels in a FITS image while remaining faithful to the FITS standard. It is written by William Pence, the principal author of the FITS standard51, and is regularly updated. Setting the definitions for all other software packages using FITS images.

Some GNU/Linux distributions have CFITSIO in their package managers, if it is available and updated, you can use it. One problem that might occur is that CFITSIO might not be configured with the --enable-reentrant option by the distribution. This option allows CFITSIO to open a file in multiple threads, it can thus provide great speed improvements. If CFITSIO was not configured with this option, any program which needs this capability will warn you and abort when you ask for multiple threads (see Multi-threaded operations).

To install CFITSIO from source, we strongly recommend that you have a look through Chapter 2 (Creating the CFITSIO library) of the CFITSIO manual and understand the options you can pass to $ ./configure (they aren’t too much). This is a very basic package for most astronomical software and it is best that you configure it nicely with your system. Once you download the source and unpack it, the following configure script should be enough for most purposes. Don’t forget to read chapter two of the manual though, for example the second option is only for 64bit systems. The manual also explains how to check if it has been installed correctly.

CFITSIO comes with two executable files called fpack and funpack. From their manual: they “are standalone programs for compressing and uncompressing images and tables that are stored in the FITS (Flexible Image Transport System) data format. They are analogous to the gzip and gunzip compression programs except that they are optimized for the types of astronomical images that are often stored in FITS format”. The commands below will compile and install them on your system along with CFITSIO. They are not essential for Gnuastro, since they are just wrappers for functions within CFITSIO, but they can come in handy. The make utils command is only available for versions above 3.39, it will build these executable files along with several other executable test files which are deleted in the following commands before the installation (otherwise the test files will also be installed).

The commands necessary to decompress, build and install CFITSIO from source are described below. Let’s assume you have downloaded cfitsio_latest.tar.gz and are in the same directory:

$ tar xf cfitsio_latest.tar.gz
$ cd cfitsio
$ ./configure --prefix=/usr/local --enable-sse2 --enable-reentrant
$ make
$ make utils
$ ./testprog > testprog.lis
$ diff testprog.lis testprog.out    # Should have no output
$ cmp testprog.std     # Should have no output
$ rm cookbook fitscopy imcopy smem speed testprog
$ sudo make install

Previous: , Up: Mandatory dependencies   [Contents][Index] WCSLIB

WCSLIB is written and maintained by one of the authors of the World Coordinate System (WCS) definition in the FITS standard52, Mark Calabretta. It might be already built and ready in your distribution’s package management system. However, here the installation from source is explained, for the advantages of installation from source please see Mandatory dependencies. To install WCSLIB you will need to have CFITSIO already installed, see CFITSIO.

WCSLIB also has plotting capabilities which use PGPLOT (a plotting library for C). If you wan to use those capabilities in WCSLIB, PGPLOT provides the PGPLOT installation instructions. However PGPLOT is old53, so its installation is not easy, there are also many great modern WCS plotting tools (mostly in written in Python). Hence, if you will not be using those plotting functions in WCSLIB, you can configure it with the --without-pgplot option as shown below.

If you have the cURL library 54 on your system and you installed CFITSIO version 3.42 or later, you will need to also link with the cURL library at configure time (through the -lcurl option as shown below). CFITSIO uses the cURL library for its HTTPS (or HTTP Secure55) support and if it is present on your system, CFITSIO will depend on it. Therefore, if ./configure command below fails (you don’t have the cURL library), then remove this option and rerun it.

Let’s assume you have downloaded wcslib.tar.bz2 and are in the same directory, to configure, build, check and install WCSLIB follow the steps below.

$ tar xf wcslib.tar.bz2

## In the `cd' command, replace `X.X' with version number.
$ cd wcslib-X.X

## If `./configure' fails, remove `-lcurl' and run again.
$ ./configure LIBS="-pthread -lcurl -lm" --without-pgplot     \
$ make
$ make check
$ sudo make install

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3.1.2 Optional dependencies

The libraries listed here are only used for very specific applications, therefore if you don’t want these operations, Gnuastro will be built and installed without them and you don’t have to have the dependencies.

If the ./configure script can’t find these requirements, it will warn you in the end that they are not present and notify you of the operation(s) you can’t do due to not having them. If the output you request from a program requires a missing library, that program is going to warn you and abort. In the case of program dependencies (like GPL GhostScript), if you install them at a later time, the program will run. This is because if required libraries are not present at build time, the executables cannot be built, but an executable is called by the built program at run time so if it becomes available, it will be used. If you do install an optional library later, you will have to rebuild Gnuastro and reinstall it for it to take effect.

GNU Libtool

Libtool is a program to simplify managing of the libraries to build an executable (a program). GNU Libtool has some added functionality compared to other implementations. If GNU Libtool isn’t present on your system at configuration time, a warning will be printed and BuildProgram won’t be built or installed. The configure script will look into your search path (PATH) for GNU Libtool through the following executable names: libtool (acceptable only if it is the GNU implementation) or glibtool. See Installation directory for more on PATH.

GNU Libtool (the binary/executable file) is a low-level program that is probably already present on your system, and if not, is available in your operating system package manager56. If you want to install GNU Libtool’s latest version from source, please visit its webpage.

Gnuastro’s tarball is shipped with an internal implementation of GNU Libtool. Even if you have GNU Libtool, Gnuastro’s internal implementation is used for the building and installation of Gnuastro. As a result, you can still build, install and use Gnuastro even if you don’t have GNU Libtool installed on your system. However, this internal Libtool does not get installed. Therefore, after Gnuastro’s installation, if you want to use BuildProgram to compile and link your own C source code which uses the Gnuastro library, you need to have GNU Libtool available on your system (independent of Gnuastro). See Review of library fundamentals to learn more about libraries.


Git is one of the most common version control systems (see Version controlled source). When libgit2 is present, and Gnuastro’s programs are run within a version controlled directory, outputs will contain the version number of the working directory’s repository for future reproducibility. See the COMMIT keyword header in Output headers for a discussion.


libjpeg is only used by ConvertType to read from and write to JPEG images, see Recognized file formats. libjpeg is a very basic library that provides tools to read and write JPEG images, most Unix-like graphic programs and libraries use it. Therefore you most probably already have it installed. libjpeg-turbo is an alternative to libjpeg. It uses Single instruction, multiple data (SIMD) instructions for ARM based systems that significantly decreases the processing time of JPEG compression and decompression algorithms.


libtiff is used by ConvertType and the libraries to read TIFF images, see Recognized file formats. libtiff is a very basic library that provides tools to read and write TIFF images, most Unix-like operating system graphic programs and libraries use it. Therefore even if you don’t have it installed, it must be easily available in your package manager.

GPL Ghostscript

GPL Ghostscript’s executable (gs) is called by ConvertType to compile a PDF file from a source PostScript file, see ConvertType. Therefore its headers (and libraries) are not needed. With a very high probability you already have it in your GNU/Linux distribution. Unfortunately it does not follow the standard GNU build style so installing it is very hard. It is best to rely on your distribution’s package managers for this.

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3.1.3 Bootstrapping dependencies

Bootstrapping is only necessary if you have decided to obtain the full version controlled history of Gnuastro, see Version controlled source and Bootstrapping. Using the version controlled source enables you to always be up to date with the most recent development work of Gnuastro (bug fixes, new functionalities, improved algorithms and etc). If you have downloaded a tarball (see Downloading the source), then you can ignore this subsection.

To successfully run the bootstrapping process, there are some additional dependencies to those discussed in the previous subsections. These are low level tools that are used by a large collection of Unix-like operating systems programs, therefore they are most probably already available in your system. If they are not already installed, you should be able to easily find them in any GNU/Linux distribution package management system (apt-get, yum, pacman and etc). The short names in parenthesis in typewriter font after the package name can be used to search for them in your package manager. For the GNU Portability Library, GNU Autoconf Archive and TeX Live, it is recommended to use the instructions here, not your operating system’s package manager.

GNU Portability Library (Gnulib)

To ensure portability for a wider range of operating systems (those that don’t include GNU C library, namely glibc), Gnuastro depends on the GNU portability library, or Gnulib. Gnulib keeps a copy of all the functions in glibc, implemented (as much as possible) to be portable to other operating systems. The bootstrap script can automatically clone Gnulib (as a gnulib/ directory inside Gnuastro), however, as described in Bootstrapping this is not recommended.

The recommended way to bootstrap Gnuastro is to first clone Gnulib and the Autoconf archives (see below) into a local directory outside of Gnuastro. Let’s call it DEVDIR57 (which you can set to any directory). Currently in Gnuastro, both Gnulib and Autoconf archives have to be cloned in the same top directory58 like the case here59:

$ DEVDIR=/home/yourname/Development
$ cd $DEVDIR
$ git clone git://
$ git clone git://

You now have the full version controlled source of these two repositories in separate directories. Both these packages are regularly updated, so every once in a while, you can run $ git pull within them to get any possible updates.

GNU Automake (automake)

GNU Automake will build the files in each sub-directory using the (hand-written) files. The Makefile.ins are subsequently used to generate the Makefiles when the user runs ./configure before building.

GNU Autoconf (autoconf)

GNU Autoconf will build the configure script using the configurations we have defined (hand-written) in

GNU Autoconf Archive

These are a large collection of tests that can be called to run at ./configure time. See the explanation under GNU Portability Library above for instructions on obtaining it and keeping it up to date.

GNU Libtool (libtool)

GNU Libtool is in charge of building all the libraries in Gnuastro. The libraries contain functions that are used by more than one program and are installed for use in other programs. They are thus put in a separate directory (lib/).

GNU help2man (help2man)

GNU help2man is used to convert the output of the --help option (--help) to the traditional Man page (Man pages).

LaTeX and some TeX packages

Some of the figures in this book are built by LaTeX (using the PGF/TikZ package). The LaTeX source for those figures is version controlled for easy maintenance not the actual figures. So the ./boostrap script will run LaTeX to build the figures. The best way to install LaTeX and all the necessary packages is through TeX live which is a package manager for TeX related tools that is independent of any operating system. It is thus preferred to the TeX Live versions distributed by your operating system.

To install TeX Live, go to the webpage and download the appropriate installer by following the “download” link. Note that by default the full package repository will be downloaded and installed (around 4 Giga Bytes) which can take very long to download and to update later. However, most packages are not needed by everyone, it is easier, faster and better to install only the “Basic scheme” (consisting of only the most basic TeX and LaTeX packages, which is less than 200 Mega bytes)60.

After the installation, be sure to set the environment variables as suggested in the end of the outputs. Any time you confront (need) a package you don’t have, simply install it with a command like below (similar to how you install software from your operating system’s package manager)61. To install all the necessary TeX packages for a successful Gnuastro bootstrap, run this command:

$ su
# tlmgr install epsf jknapltx caption biblatex biber iftex \
                etoolbox logreq xstring xkeyval pgf ms     \
                xcolor pgfplots times rsfs pstools epspdf
ImageMagick (imagemagick)

ImageMagick is a wonderful and robust program for image manipulation on the command-line. bootsrap uses it to convert the book images into the formats necessary for the various book formats.

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3.1.4 Dependencies from package managers

The most basic way to install a package on your system is to build the packages from source yourself. Alternatively, you can use your operating system’s package manager to download pre-compiled files and install them. The latter choice is easier and faster. However, we recommend that you build the Mandatory dependencies yourself from source (all necessary commands and links are given in the respective section). Here are some basic reasons behind this recommendation.

  1. Your distribution’s pre-built package might not be the most recent release.
  2. For each package, Gnuastro might preform better (or require) certain configuration options that your distribution’s package managers didn’t add for you. If present, these configuration options are explained during the installation of each in the sections below (for example in CFITSIO). When the proper configuration has not been set, the programs should complain and inform you.
  3. For the libraries, they might separate the binary file from the header files which can cause confusion, see Known issues.
  4. Like any other tool, the science you derive from Gnuastro’s tools highly depend on these lower level dependencies, so generally it is much better to have a close connection with them. By reading their manuals, installing them and staying up to date with changes/bugs in them, your scientific results and understanding (of what is going on, and thus how you interpret your scientific results) will also correspondingly improve.

Based on your package manager, you can use any of the following commands to install the mandatory and optional dependencies. If your package manager isn’t included in the list below, please send us the respective command, so we add it. Gnuastro itself if also already packaged in some package managers (for example Debian or Homebrew).

As discussed above, we recommend installing the mandatory dependencies manually from source (see Mandatory dependencies). Therefore, in each command below, first the optional dependencies are given. The mandatory dependencies are included after an empty line. If you would also like to install the mandatory dependencies with your package manager, just ignore the empty line.

For better archivability and compression ratios, Gnuastro’s recommended tarball compression format is with the Lzip program, see Release tarball. Therefore, the package manager commands below also contain Lzip.

apt-get (Debian-based OSs: Debian, Ubuntu, Linux Mint, and etc)

Debian is one of the oldest GNU/Linux distributions62. It thus has a very extended user community and a robust internal structure and standards. All of it is free software and based on the work of volunteers around the world. Many distributions are thus derived from it, for example Ubuntu and Linux Mint. This arguably makes Debian-based OSs the largest, and most used, class of GNU/Linux distributions. All of them use Debian’s Advanced Packaging Tool (APT, for example apt-get) for managing packages.

$ sudo apt-get install ghostscript libtool-bin libjpeg-dev  \
                       libtiff-dev libgit2-dev lzip         \
                       libgsl0-dev libcfitsio-dev wcslib-dev
yum (Red Hat-based OSs: Red Hat, Fedora, CentOS, Scientific Linux, and etc)

Red Hat Enterprise Linux (RHEL) is released by Red Hat Inc. RHEL requires paid subscriptions for use of its binaries and support. But since it is free software, many other teams use its code to spin-off their own distributions based on RHEL. Red Hat-based GNU/Linux distributions initially used the “Yellowdog Updated, Modifier” (YUM) package manager, which has been replaced by “Dandified yum” (DNF). If the latter isn’t available on your system, you can use yum instead of dnf in the command below.

$ sudo dnf install ghostscript libtool libjpeg-devel        \
                   libtiff-devel libgit2-devel lzip         \
                   gsl-devel cfitsio-devel wcslib-devel
brew (macOS)

macOS is the operating system used on Apple devices. macOS does not come with a package manager pre-installed, but several widely used, third-party package managers exist, such as Homebrew or MacPorts. Both are free software. Currently we have only tested Gnuastro’s installation with Homebrew as described below.

If not already installed, first obtain Homebrew by following the instructions at Homebrew manages packages in different ‘taps’. To install WCSLIB (discussed in Mandatory dependencies) via Homebrew you will need to tap into brewsci/science first (the tap may change in the future, but can be found by calling brew search wcslib).

$ brew install ghostscript libtool libjpeg libtiff          \
               libgit2 lzip                                 \
               gsl cfitsio
$ brew tap brewsci/science
$ brew install wcslib
pacman (Arch Linux)

Arch Linux is a smaller GNU/Linux distribution, which follows the KISS principle (“keep it simple, stupid”) as a general guideline. It “focuses on elegance, code correctness, minimalism and simplicity, and expects the user to be willing to make some effort to understand the system’s operation”. Arch Linux uses “Package manager” (Pacman) to manage its packages/components.

$ sudo pacman -S ghostscript libtool libjpeg libtiff        \
                 libgit2 lzip                               \
                 gsl cfitsio wcslib

Usually, when libraries are installed by operating system package managers, there should be no problems when configuring and building other programs from source (that depend on the libraries: Gnuastro in this case). However, in some special conditions, problems may pop-up during the configuration, building, or checking/running any of Gnuastro’s programs. The most common of such problems and their solution are discussed below.

Not finding library during configuration: If a library is installed, but during Gnuastro’s configure step the library isn’t found, then configure Gnuastro like the command below (correcting /path/to/lib). For more, see Known issues and Installation directory.

$ ./configure LDFLAGS="-I/path/to/lib"

Not finding header (.h) files while building: If a library is installed, but during Gnuastro’s make step, the library’s header (file with a .h suffix) isn’t found, then configure Gnuastro like the command below (correcting /path/to/include). For more, see Known issues and Installation directory.

$ ./configure CPPFLAGS="-I/path/to/include"

Gnuastro’s programs don’t run during check or after install: If a library is installed, but the programs don’t run due to linking problems, set the LD_LIBRARY_PATH variable like below (assuming Gnuastro is installed in /path/to/installed). For more, see Known issues and Installation directory.

$ export LD_LIBRARY_PATH="$LD_LIBRARY_PATH:/path/to/installed/lib"

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3.2 Downloading the source

Gnuastro’s source code can be downloaded in two ways. As a tarball, ready to be configured and installed on your system (as described in Quick start), see Release tarball. If you want official releases of stable versions this is the best, easiest and most common option. Alternatively, you can clone the version controlled history of Gnuastro, run one extra bootstrapping step and then follow the same steps as the tarball. This will give you access to all the most recent work that will be included in the next release along with the full project history. The process is thoroughly introduced in Version controlled source.

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3.2.1 Release tarball

A release tarball (commonly compressed) is the most common way of obtaining free and open source software. A tarball is a snapshot of one particular moment in the Gnuastro development history along with all the necessary files to configure, build, and install Gnuastro easily (see Quick start). It is very straightforward and needs the least set of dependencies (see Mandatory dependencies). Gnuastro has tarballs for official stable releases and pre-releases for testing. See Version numbering for more on the two types of releases and the formats of the version numbers. The URLs for each type of release are given below.

Official stable releases (

This URL hosts the official stable releases of Gnuastro. Always use the most recent version (see Version numbering). By clicking on the “Last modified” title of the second column, the files will be sorted by their date which you can also use to find the latest version. It is recommended to use a mirror to download these tarballs, please visit and see below.

Pre-release tar-balls (

This URL contains unofficial pre-release versions of Gnuastro. The pre-release versions of Gnuastro here are for enthusiasts to try out before an official release. If there are problems, or bugs then the testers will inform the developers to fix before the next official release. See Version numbering to understand how the version numbers here are formatted. If you want to remain even more up-to-date with the developing activities, please clone the version controlled source as described in Version controlled source.

Gnuastro’s official/stable tarball is released with two formats: Gzip (with suffix .tar.gz) and Lzip (with suffix .tar.lz). The pre-release tarballs (after version 0.3) are released only as an Lzip tarball. Gzip is a very well-known and widely used compression program created by GNU and available in most systems. However, Lzip provides a better compression ratio and more robust archival capacity. For example Gnuastro 0.3’s tarball was 2.9MB and 4.3MB with Lzip and Gzip respectively, see the Lzip webpage for more. Lzip might not be pre-installed in your operating system, if so, installing it from your operating system’s package manager or from source is very easy and fast (it is a very small program).

The GNU FTP server is mirrored (has backups) in various locations on the globe ( You can use the closest mirror to your location for a more faster download. Note that only some mirrors keep track of the pre-release (alpha) tarballs. Also note that if you want to download immediately after and announcement (see Announcements), the mirrors might need some time to synchronize with the main GNU FTP server.

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3.2.2 Version controlled source

The publicly distributed Gnuastro tar-ball (for example gnuastro-X.X.tar.gz) does not contain the revision history, it is only a snapshot of the source code at one significant instant of Gnuastro’s history (specified by the version number, see Version numbering), ready to be configured and built. To be able to develop successfully, the revision history of the code can be very useful to track when something was added or changed, also some updates that are not yet officially released might be in it.

We use Git for the version control of Gnuastro. For those who are not familiar with it, we recommend the ProGit book. The whole book is publicly available for online reading and downloading and does a wonderful job at explaining the concepts and best practices.

Let’s assume you want to keep Gnuastro in the TOPGNUASTRO directory (can be any directory, change the value below). The full version controlled history of Gnuastro can be cloned in TOPGNUASTRO/gnuastro by running the following commands63:

$ TOPGNUASTRO=/home/yourname/Research/projects/
$ git clone git://

The $TOPGNUASTRO/gnuastro directory will contain hand-written (version controlled) source code for Gnuastro’s programs, libraries, this book and the tests. All are divided into sub-directories with standard and very descriptive names. The version controlled files in the top cloned directory are either mainly in capital letters (for example THANKS and README) or mainly written in small-caps (for example and The former are non-programming, standard writing for human readers containing high-level information about the whole package. The latter are instructions to customize the GNU build system for Gnuastro. For more on Gnuastro’s source code structure, please see Developing. We won’t go any deeper here.

The cloned Gnuastro source cannot immediately be configured, compiled, or installed since it only contains hand-written files, not automatically generated or imported files which do all the hard work of the build process. See Bootstrapping for the process of generating and importing those files (its not too hard!). Once you have bootstrapped Gnuastro, you can run the standard procedures (in Quick start). Very soon after you have cloned it, Gnuastro’s main master branch will be updated on the main repository (since the developers are actively working on Gnuastro), for the best practices in keeping your local history in sync with the main repository see Synchronizing.

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The version controlled source code lacks the source files that we have not written or are automatically built. These automatically generated files are included in the distributed tar ball for each distribution (for example gnuastro-X.X.tar.gz, see Version numbering) and make it easy to immediately configure, build, and install Gnuastro. However from the perspective of version control, they are just bloatware and sources of confusion (since they are not changed by Gnuastro developers).

The process of automatically building and importing necessary files into the cloned directory is known as bootstrapping. All the instructions for an automatic bootstrapping are available in bootstrap and configured using bootstrap.conf. bootstrap and COPYING (which contains the software copyright notice) are the only files not written by Gnuastro developers but under version control to enable simple bootstrapping and legal information on usage immediately after cloning. bootstrap.conf is maintained by the GNU Portability Library (Gnulib) and this file is an identical copy, so do not make any changes in this file since it will be replaced when Gnulib releases an update. Make all your changes in bootstrap.conf.

The bootstrapping process has its own separate set of dependencies, the full list is given in Bootstrapping dependencies. They are generally very low-level and used by a very large set of commonly used programs, so they are probably already installed on your system. The simplest way to bootstrap Gnuastro is to simply run the bootstrap script within your cloned Gnuastro directory as shown below. However, please read the next paragraph before doing so (see Version controlled source for TOPGNUASTRO).

$ cd TOPGNUASTRO/gnuastro
$ ./bootstrap                      # Requires internet connection

Without any options, bootstrap will clone Gnulib within your cloned Gnuastro directory (TOPGNUASTRO/gnuastro/gnulib) and download the necessary Autoconf archives macros. So if you run bootstrap like this, you will need an internet connection every time you decide to bootstrap. Also, Gnulib is a large package and cloning it can be slow. It will also keep the full Gnulib repository within your Gnuastro repository, so if another one of your projects also needs Gnulib, and you insist on running bootstrap like this, you will have two copies. In case you regularly backup your important files, Gnulib will also slow down the backup process. Therefore while the simple invocation above can be used with no problem, it is not recommended. To do better, see the next paragraph.

The recommended way to get these two packages is thoroughly discussed in Bootstrapping dependencies (in short: clone them in the separate DEVDIR/ directory). The following commands will take you into the cloned Gnuastro directory and run the bootstrap script, while telling it to copy some files (instead of making symbolic links, with the --copy option, this is not mandatory64) and where to look for Gnulib (with the --gnulib-srcdir option). Please note that the address given to --gnulib-srcdir has to be an absolute address (so don’t use ~ or ../ for example).

$ cd $TOPGNUASTRO/gnuastro
$ ./bootstrap --copy --gnulib-srcdir=$DEVDIR/gnulib

Since Gnulib and Autoconf archives are now available in your local directories, you don’t need an internet connection every time you decide to remove all untracked files and redo the bootstrap (see box below). You can also use the same command on any other project that uses Gnulib. All the necessary GNU C library functions, Autoconf macros and Automake inputs are now available along with the book figures. The standard GNU build system (Quick start) will do the rest of the job.

Undoing the bootstrap: During the development, it might happen that you want to remove all the automatically generated and imported files. In other words, you might want to reverse the bootstrap process. Fortunately Git has a good program for this job: git clean. Run the following command and every file that is not version controlled will be removed.

git clean -fxd

It is best to commit any recent change before running this command. You might have created new files since the last commit and if they haven’t been committed, they will all be gone forever (using rm). To get a list of the non-version controlled files instead of deleting them, add the n option to git clean, so it becomes -fxdn.

Besides the bootstrap and bootstrap.conf, the bootstrapped/ directory and README-hacking file are also related to the bootstrapping process. The former hosts all the imported (bootstrapped) directories. Thus, in the version controlled source, it only contains a REAME file, but in the distributed tar-ball it also contains sub-directories filled with all bootstrapped files. README-hacking contains a summary of the bootstrapping process discussed in this section. It is a necessary reference when you haven’t built this book yet. It is thus not distributed in the Gnuastro tarball.

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The bootstrapping script (see Bootstrapping) is not regularly needed: you mainly need it after you have cloned Gnuastro (once) and whenever you want to re-import the files from Gnulib, or Autoconf archives65 (not too common). However, Gnuastro developers are constantly working on Gnuastro and are pushing their changes to the official repository. Therefore, your local Gnuastro clone will soon be out-dated. Gnuastro has two mailing lists dedicated to its developing activities (see Developing mailing lists). Subscribing to them can help you decide when to synchronize with the official repository.

To pull all the most recent work in Gnuastro, run the following command from the top Gnuastro directory. If you don’t already have a built system, ignore make distclean. The separate steps are described in detail afterwards.

$ make distclean && git pull && autoreconf -f

You can also run the commands separately:

$ make distclean
$ git pull
$ autoreconf -f

If Gnuastro was already built in this directory, you don’t want some outputs from the previous version being mixed with outputs from the newly pulled work. Therefore, the first step is to clean/delete all the built files with make distclean. Fortunately the GNU build system allows the separation of source and built files (in separate directories). This is a great feature to keep your source directory clean and you can use it to avoid the cleaning step. Gnuastro comes with a script with some useful options for this job. It is useful if you regularly pull recent changes, see Separate build and source directories.

After the pull, we must re-configure Gnuastro with autoreconf -f (part of GNU Autoconf). It will update the ./configure script and all the Makefile.in66 files based on the hand-written configurations (in and the files). After running autoreconf -f, a warning about TEXI2DVI might show up, you can ignore that.

The most important reason for re-building Gnuastro’s build system is to generate/update the version number for your updated Gnuastro snapshot. This generated version number will include the commit information (see Version numbering). The version number is included in nearly all outputs of Gnuastro’s programs, therefore it is vital for reproducing an old result.

As a summary, be sure to run ‘autoreconf -f’ after every change in the Git history. This includes synchronization with the main server or even a commit you have made yourself.

If you would like to see what has changed since you last synchronized your local clone, you can take the following steps instead of the simple command above (don’t type anything after #):

$ git checkout master             # Confirm if you are on master.
$ git fetch origin                # Fetch all new commits from server.
$ git log master..origin/master   # See all the new commit messages.
$ git merge origin/master         # Update your master branch.
$ autoreconf -f                   # Update the build system.

By default git log prints the most recent commit first, add the --reverse option to see the changes chronologically. To see exactly what has been changed in the source code along with the commit message, add a -p option to the git log.

If you want to make changes in the code, have a look at Developing to get started easily. Be sure to commit your changes in a separate branch (keep your master branch to follow the official repository) and re-run autoreconf -f after the commit. If you intend to send your work to us, you can safely use your commit since it will be ultimately recorded in Gnuastro’s official history. If not, please upload your separate branch to a public hosting service, for example GitLab, and link to it in your report/paper. Alternatively, run make distcheck and upload the output gnuastro-X.X.X.XXXX.tar.gz to a publicly accessible webpage so your results can be considered scientific (reproducible) later.

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3.3 Build and install

This section is basically a longer explanation to the sequence of commands given in Quick start. If you didn’t have any problems during the Quick start steps, you want to have all the programs of Gnuastro installed in your system, you don’t want to change the executable names during or after installation, you have root access to install the programs in the default system wide directory, the Letter paper size of the print book is fine for you or as a summary you don’t feel like going into the details when everything is working, you can safely skip this section.

If you have any of the above problems or you want to understand the details for a better control over your build and install, read along. The dependencies which you will need prior to configuring, building and installing Gnuastro are explained in Dependencies. The first three steps in Quick start need no extra explanation, so we will skip them and start with an explanation of Gnuastro specific configuration options and a discussion on the installation directory in Configuring, followed by some smaller subsections: Tests, A4 print book, and Known issues which explains the solutions to known problems you might encounter in the installation steps and ways you can solve them.

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3.3.1 Configuring

The $ ./configure step is the most important step in the build and install process. All the required packages, libraries, headers and environment variables are checked in this step. The behaviors of make and make install can also be set through command line options to this command.

The configure script accepts various arguments and options which enable the final user to highly customize whatever she is building. The options to configure are generally very similar to normal program options explained in Arguments and options. Similar to all GNU programs, you can get a full list of the options along with a short explanation by running

$ ./configure --help

A complete explanation is also included in the INSTALL file. Note that this file was written by the authors of GNU Autoconf (which builds the configure script), therefore it is common for all programs which use the $ ./configure script for building and installing, not just Gnuastro. Here we only discuss cases where you don’t have super-user access to the system and if you want to change the executable names. But before that, a review of the options to configure that are particular to Gnuastro are discussed.

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Most of the options to configure (which are to do with building) are similar for every program which uses this script. Here the options that are particular to Gnuastro are discussed. The next topics explain the usage of other configure options which can be applied to any program using the GNU build system (through the configure script).


Compile/build Gnuastro with debugging information, no optimization and without shared libraries.

In order to allow more efficient programs when using Gnuastro (after the installation), by default Gnuastro is built with a 3rd level (a very high level) optimization and no debugging information. By default, libraries are also built for static and shared linking (see Linking). However, when there are crashes or unexpected behavior, these three features can hinder the process of localizing the problem. This configuration option is identical to manually calling the configuration script with CFLAGS="-g -O0" --disable-shared.

In the (rare) situations where you need to do your debugging on the shared libraries, don’t use this option. Instead run the configure script by explicitly setting CFLAGS like this:

$ ./configure CFLAGS="-g -O0"

Do the make check tests through Valgrind. Therefore, if any crashes or memory-related issues (segmentation faults in particular) occur in the tests, the output of Valgrind will also be put in the tests/test-suite.log file without having to manually modify the check scripts. This option will also activate Gnuastro’s debug mode (see the --enable-debug configure-time option described above).

Valgrind is free software. It is a program for easy checking of memory-related issues in programs. It runs a program within its own controlled environment and can thus identify the exact line-number in the program’s source where a memory-related issue occurs. However, it can significantly slow-down the tests. So this option is only useful when a segmentation fault is found during make check.


Only build and install progname along with any other program that is enabled in this fashion. progname is the name of the executable without the ast, for example crop for Crop (with the executable name of astcrop). If this option is called for any of the programs in Gnuastro, any program which is not explicitly enabled will not be built or installed.


Do not build or install the program named progname. This is very similar to the --enable-progname, but will build and install all the other programs except this one.


Enable checks on the GNU Portability Library (Gnulib). Gnulib is used by Gnuastro to enable users of non-GNU based operating systems (that don’t use GNU C library or glibc) to compile and use the advanced features that this library provides. We make extensive use of such functions. If you give this option to $ ./configure, when you run $ make check, first the functions in Gnulib will be tested, then the Gnuastro executables. If your operating system does not support glibc or has an older version of it and you have problems in the build process ($ make), you can give this flag to configure to see if the problem is caused by Gnulib not supporting your operating system or Gnuastro, see Known issues.


Do not print a guiding message during the GNU Build process of Quick start. By default, after each step, a message is printed guiding the user what the next command should be. Therefore, after ./configure, it will suggest running make. After make, it will suggest running make check and so on. If Gnuastro is configured with this option, for example

$ ./configure --disable-guide-message

Then these messages will not be printed after any step (like most programs). For people who are not yet fully accustomed to this build system, these guidelines can be very useful and encouraging. However, if you find those messages annoying, use this option.

Note: If some programs are enabled and some are disabled, it is equivalent to simply enabling those that were enabled. Listing the disabled programs is redundant.

The tests of some programs might depend on the outputs of the tests of other programs. For example MakeProfiles is one the first programs to be tested when you run $ make check. MakeProfiles’ test outputs (FITS images) are inputs to many other programs (which in turn provide inputs for other programs). Therefore, if you don’t install MakeProfiles for example, the tests for many the other programs will be skipped. To avoid this, in one run, you can install all the programs and run the tests but not install. If everything is working correctly, you can run configure again with only the programs you want. However, don’t run the tests and directly install after building.

Next: , Previous: , Up: Configuring   [Contents][Index] Installation directory

One of the most commonly used options to ./configure is --prefix, it is used to define the directory that will host all the installed files (or the “prefix” in their final absolute file name). For example, when you are using a server and you don’t have administrator or root access. In this example scenario, if you don’t use the --prefix option, you won’t be able to install the built files and thus access them from anywhere without having to worry about where they are installed. However, once you prepare your startup file to look into the proper place (as discussed thoroughly below), you will be able to easily use this option and benefit from any software you want to install without having to ask the system administrators or install and use a different version of a software that is already installed on the server.

The most basic way to run an executable is to explicitly write its full file name (including all the directory information) and run it. One example is running the configuration script with the $ ./configure command (see Quick start). By giving a specific directory (the current directory or ./), we are explicitly telling the shell to look in the current directory for an executable file named ‘configure’. Directly specifying the directory is thus useful for executables in the current (or nearby) directories. However, when the program (an executable file) is to be used a lot, specifying all those directories will become a significant burden. For example, the ls executable lists the contents in a given directory and it is (usually) installed in the /usr/bin/ directory by the operating system maintainers. Therefore, if using the full address was the only way to access an executable, each time you wanted a listing of a directory, you would have to run the following command (which is very inconvenient, both in writing and in remembering the various directories).

$ /usr/bin/ls

To address this problem, we have the PATH environment variable. To understand it better, we will start with a short introduction to the shell variables. Shell variable values are basically treated as strings of characters. For example, it doesn’t matter if the value is a name (string of alphabetic characters), or a number (string of numeric characters), or both. You can define a variable and a value for it by running

$ myvariable1=a_test_value
$ myvariable2="a test value"

As you see above, if the value contains white space characters, you have to put the whole value (including white space characters) in double quotes ("). You can see the value it represents by running

$ echo $myvariable1
$ echo $myvariable2

If a variable has no value or it wasn’t defined, the last command will only print an empty line. A variable defined like this will be known as long as this shell or terminal is running. Other terminals will have no idea it existed. The main advantage of shell variables is that if they are exported67, subsequent programs that are run within that shell can access their value. So by changing their value, you can change the “environment” of a program which uses them. The shell variables which are accessed by programs are therefore known as “environment variables”68. You can see the full list of exported variables that your shell recognizes by running:

$ printenv

HOME is one commonly used environment variable, it is any user’s (the one that is logged in) top directory. Try finding it in the command above. It is used so often that the shell has a special expansion (alternative) for it: ‘~’. Whenever you see file names starting with the tilde sign, it actually represents the value to the HOME environment variable, so ~/doc is the same as $HOME/doc.

Another one of the most commonly used environment variables is PATH, it is a list of directories to search for executable names. Its value is a list of directories (separated by a colon, or ‘:’). When the address of the executable is not explicitly given (like ./configure above), the system will look for the executable in the directories specified by PATH. If you have a computer nearby, try running the following command to see which directories your system will look into when it is searching for executable (binary) files, one example is printed here (notice how /usr/bin, in the ls example above, is one of the directories in PATH):

$ echo $PATH

By default PATH usually contains system-wide directories, which are readable (but not writable) by all users, like the above example. Therefore if you don’t have root (or administrator) access, you need to add another directory to PATH which you actually have write access to. The standard directory where you can keep installed files (not just executables) for your own user is the ~/.local/ directory. The names of hidden files start with a ‘.’ (dot), so it will not show up in your common command-line listings, or on the graphical user interface. You can use any other directory, but this is the most recognized.

The top installation directory will be used to keep all the package’s components: programs (executables), libraries, include (header) files, shared data (like manuals), or configuration files (see Review of library fundamentals for a thorough introduction to headers and linking). So it commonly has some of the following sub-directories for each class of installed components respectively: bin/, lib/, include/ man/, share/, etc/. Since the PATH variable is only used for executables, you can add the ~/.local/bin directory (which keeps the executables/programs or more generally, “binary” files) to PATH with the following command. As defined below, first the existing value of PATH is used, then your given directory is added to its end and the combined value is put back in PATH (run ‘$ echo $PATH’ afterwards to check if it was added).

$ PATH=$PATH:~/.local/bin

Any executable that you installed in ~/.local/bin will now be usable without having to remember and write its full address. However, as soon as you leave/close your current terminal session, this modified PATH variable will be forgotten. Adding the directories which contain executables to the PATH environment variable each time you start a terminal is also very inconvenient and prone to errors. Fortunately, there are standard ‘startup files’ defined by your shell precisely for this (and other) purposes. There is a special startup file for every significant starting step:

/etc/profile and everything in /etc/profile.d/

These startup scripts are called when your whole system starts (for example after you turn on your computer). Therefore you need administrator or root privileges to access or modify them.


If you are using (GNU) Bash as your shell, the commands in this file are run, when you log in to your account through Bash. Most commonly when you login through the virtual console (where there is no graphic user interface).


If you are using (GNU) Bash as your shell, the commands here will be run each time you start a terminal and are already logged in. For example, when you open your terminal emulator in the graphic user interface.

For security reasons, it is highly recommended to directly type in your HOME directory value by hand in startup files instead of using variables. So in the following, let’s assume your user name is ‘name’ (so ~ may be replaced with /home/name). To add ~/.local/bin to your PATH automatically on any startup file, you have to “export” the new value of PATH in the startup file that is most relevant to you by adding this line:

export PATH=$PATH:/home/name/.local/bin

Now that you know your system will look into ~/.local/bin for executables, you can tell Gnuastro’s configure script to install everything in the top ~/.local directory using the --prefix option. When you subsequently run $ make install, all the install-able files will be put in their respective directory under ~/.local/ (the executables in ~/.local/bin, the compiled library files in ~/.local/lib, the library header files in ~/.local/include and so on, to learn more about these different files, please see Review of library fundamentals). Note that tilde (‘~’) expansion will not happen if you put a ‘=’ between --prefix and ~/.local69, so we have avoided the = character here which is optional in GNU-style options, see Options.

$ ./configure --prefix ~/.local

You can install everything (including libraries like GSL, CFITSIO, or WCSLIB which are Gnuastro’s mandatory dependencies, see Mandatory dependencies) locally by configuring them as above. However, recall that PATH is only for executable files, not libraries and that libraries can also depend on other libraries. For example WCSLIB depends on CFITSIO and Gnuastro needs both. Therefore, when you installed a library in a non-recognized directory, you have to guide the program that depends on them to look into the necessary library and header file directories. To do that, you have to define the LDFLAGS and CPPFLAGS environment variables respectively. This can be done while calling ./configure as shown below:

$ ./configure LDFLAGS=-L/home/name/.local/lib            \
              CPPFLAGS=-I/home/name/.local/include       \
              --prefix ~/.local

It can be annoying/buggy to do this when configuring every software that depends on such libraries. Hence, you can define these two variables in the most relevant startup file (discussed above). The convention on using these variables doesn’t include a colon to separate values (as PATH-like variables do), they use white space characters and each value is prefixed with a compiler option70: note the -L and -I above (see Options), for -I see Headers, and for -L, see Linking. Therefore we have to keep the value in double quotation signs to keep the white space characters and adding the following two lines to the startup file of choice:

export LDFLAGS="$LDFLAGS -L/home/name/.local/lib"
export CPPFLAGS="$CPPFLAGS -I/home/name/.local/include"

Dynamic libraries are linked to the executable every time you run a program that depends on them (see Linking to fully understand this important concept). Hence dynamic libraries also require a special path variable called LD_LIBRARY_PATH (same formatting as PATH). To use programs that depend on these libraries, you need to add ~/.local/lib to your LD_LIBRARY_PATH environment variable by adding the following line to the relevant start-up file:

export LD_LIBRARY_PATH=$LD_LIBRARY_PATH:/home/name/.local/lib

If you also want to access the Info (see Info) and man pages (see Man pages) documentations add ~/.local/share/info and ~/.local/share/man to your INFOPATH71 and MANPATH environment variables respectively.

A final note is that order matters in the directories that are searched for all the variables discussed above. In the examples above, the new directory was added after the system specified directories. So if the program, library or manuals are found in the system wide directories, the user directory is no longer searched. If you want to search your local installation first, put the new directory before the already existing list, like the example below.

export LD_LIBRARY_PATH=/home/name/.local/lib:$LD_LIBRARY_PATH

This is good when a library, for example CFITSIO, is already present on the system, but the system-wide install wasn’t configured with the correct configuration flags (see CFITSIO), or you want to use a newer version and you don’t have administrator or root access to update it on the whole system/server. If you update LD_LIBRARY_PATH by placing ~/.local/lib first (like above), the linker will first find the CFITSIO you installed for yourself and link with it. It thus will never reach the system-wide installation.

There are important security problems with using local installations first: all important system-wide executables and libraries (important executables like ls and cp, or libraries like the C library) can be replaced by non-secure versions with the same file names and put in the customized directory (~/.local in this example). So if you choose to search in your customized directory first, please be sure to keep it clean from executables or libraries with the same names as important system programs or libraries.

Summary: When you are using a server which doesn’t give you administrator/root access AND you would like to give priority to your own built programs and libraries, not the version that is (possibly already) present on the server, add these lines to your startup file. See above for which startup file is best for your case and for a detailed explanation on each. Don’t forget to replace ‘/YOUR-HOME-DIR’ with your home directory (for example ‘/home/your-id’):

export PATH="/YOUR-HOME-DIR/.local/bin:$PATH"
export LDFLAGS="-L/YOUR-HOME-DIR/.local/lib $LDFLAGS"
export MANPATH="/YOUR-HOME-DIR/.local/share/man/:$MANPATH"
export CPPFLAGS="-I/YOUR-HOME-DIR/.local/include $CPPFLAGS"
export INFOPATH="/YOUR-HOME-DIR/.local/share/info/:$INFOPATH"

Afterwards, you just need to add an extra --prefix=/YOUR-HOME-DIR/.local to the ./configure command of the software that you intend to install. Everything else will be the same as a standard build and install, see Quick start.

Next: , Previous: , Up: Configuring   [Contents][Index] Executable names

At first sight, the names of the executables for each program might seem to be uncommonly long, for example astnoisechisel or astcrop. We could have chosen terse (and cryptic) names like most programs do. We chose this complete naming convention (something like the commands in TeX) so you don’t have to spend too much time remembering what the name of a specific program was. Such complete names also enable you to easily search for the programs.

To facilitate typing the names in, we suggest using the shell auto-complete. With this facility you can find the executable you want very easily. It is very similar to file name completion in the shell. For example, simply by typing the letters below (where [TAB] stands for the Tab key on your keyboard)

$ ast[TAB][TAB]

you will get the list of all the available executables that start with ast in your PATH environment variable directories. So, all the Gnuastro executables installed on your system will be listed. Typing the next letter for the specific program you want along with a Tab, will limit this list until you get to your desired program.

In case all of this does not convince you and you still want to type short names, some suggestions are given below. You should have in mind though, that if you are writing a shell script that you might want to pass on to others, it is best to use the standard name because other users might not have adopted the same customization. The long names also serve as a form of documentation in such scripts. A similar reasoning can be given for option names in scripts: it is good practice to always use the long formats of the options in shell scripts, see Options.

The simplest solution is making a symbolic link to the actual executable. For example let’s assume you want to type ic to run Crop instead of astcrop. Assuming you installed Gnuastro executables in /usr/local/bin (default) you can do this simply by running the following command as root:

# ln -s /usr/local/bin/astcrop /usr/local/bin/ic

In case you update Gnuastro and a new version of Crop is installed, the default executable name is the same, so your custom symbolic link still works.

The installed executable names can also be set using options to $ ./configure, see Configuring. GNU Autoconf (which configures Gnuastro for your particular system), allows the builder to change the name of programs with the three options --program-prefix, --program-suffix and --program-transform-name. The first two are for adding a fixed prefix or suffix to all the programs that will be installed. This will actually make all the names longer! You can use it to add versions of program names to the programs in order to simultaneously have two executable versions of a program.

The third configure option allows you to set the executable name at install time using the SED program. SED is a very useful ‘stream editor’. There are various resources on the internet to use it effectively. However, we should caution that using configure options will change the actual executable name of the installed program and on every re-install (an update for example), you have to also add this option to keep the old executable name updated. Also note that the documentation or configuration files do not change from their standard names either.

For example, let’s assume that typing ast on every invocation of every program is really annoying you! You can remove this prefix from all the executables at configure time by adding this option:

$ ./configure --program-transform-name='s/ast/ /'

Previous: , Up: Configuring   [Contents][Index] Configure and build in RAM

Gnuastro’s configure and build process (the GNU build system) involves the creation, reading, and modification of a large number of files (input/output, or I/O). Therefore file I/O issues can directly affect the work of developers who need to configure and build Gnuastro numerous times. Some of these issues are listed below:

One solution to address both these problems is to use the tmpfs file system. Any file in tmpfs is actually stored in the RAM (and possibly SAWP), not on HDDs or SSDs. The RAM is built for extensive and fast I/O. Therefore the large number of file I/Os associated with configuring and building will not harm the HDDs or SSDs. Due to the volatile nature of RAM, files in the tmpfs file-system will be permanently lost after a power-off. Since all configured and built files are derivative files (not files that have been directly written by hand) there is no problem in this and this feature can be considered as an automatic cleanup.

The modern GNU C library (and thus the Linux kernel) defines the /dev/shm directory for this purpose in the RAM (POSIX shared memory). To build in it, you can use the GNU build system’s ability to build in a separate directory (not necessarily in the source directory) as shown below. Just set SRCDIR as the address of Gnuastro’s top source directory (for example, the unpacked tarball).

$ mkdir /dev/shm/tmp-gnuastro-build
$ cd /dev/shm/tmp-gnuastro-build
$ SRCDIR/configure --srcdir=SRCDIR
$ make

Gnuastro comes with a script to simplify this process of configuring and building in a different directory (a “clean” build), for more see Separate build and source directories.

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3.3.2 Separate build and source directories

The simple steps of Quick start will mix the source and built files. This can cause inconvenience for developers or enthusiasts following the the most recent work (see Version controlled source). The current section is mainly focused on this later group of Gnuastro users. If you just install Gnuastro on major releases (following Announcements), you can safely ignore this section.

When it is necessary to keep the source (which is under version control), but not the derivative (built) files (after checking or installing), the best solution is to keep the source and the built files in separate directories. One application of this is already discussed in Configure and build in RAM.

To facilitate this process of configuring and building in a separate directory, Gnuastro comes with the developer-build script. It is available in the top source directory and is not installed. It will make a directory under a given top-level directory (given to --top-build-dir) and build Gnuastro in there directory. It thus keeps the source completely separated from the built files. For easy access to the built files, it also makes a symbolic link to the built directory in the top source files called build.

When run without any options, default values will be used for its configuration. As with Gnuastro’s programs, you can inspect the default values with -P (or --printparams, the output just looks a little different here). The default top-level build directory is /dev/shm: the shared memory directory in RAM on GNU/Linux systems as described in Configure and build in RAM.

Besides these, it also has some features to facilitate the job of developers or bleeding edge users like the --debug option to do a fast build, with debug information, no optimization, and no shared libraries. Here is the full list of options you can feed to this script to configure its operations.

Not all Gnuastro’s common program behavior usable here: developer-build is just a non-installed script with a very limited scope as described above. It thus doesn’t have all the common option behaviors or configuration files for example.

White space between option and value: developer-build doesn’t accept an = sign between the options and their values. It also needs at least one character between the option and its value. Therefore -n 4 or --numthreads 4 are acceptable, while -n4, -n=4, or --numthreads=4 aren’t. Finally multiple short option names cannot be merged: for example you can say -c -n 4, but unlike Gnuastro’s programs, -cn4 is not acceptable.

Reusable for other packages: This script can be used in any software which is configured and built using the GNU Build System. Just copy it in the top source directory of that software and run it from there.

-b STR
--top-build-dir STR

The top build directory to make a directory for the build. If this option isn’t called, the top build directory is /dev/shm (only available in GNU/Linux operating systems, see Configure and build in RAM).


Print the version string of Gnuastro that will be used in the build. This string will be appended to the directory name containing the built files.


Run autoreconf -f before building the package. In Gnuastro, this is necessary when a new commit has been made to the project history. In Gnuastro’s build system, the Git description will be used as the version, see Version numbering and Synchronizing.


Delete the contents of the build directory (clean it) before starting the configuration and building of this run.

This is useful when you have recently pulled changes from the main Git repository, or committed a change your self and ran autoreconf -f, see Synchronizing. After running GNU Autoconf, the version will be updated and you need to do a clean build.


Build with debugging flags (for example to use in GNU Debugger, also known as GDB, or Valgrind), disable optimization and also the building of shared libraries. Similar to running the configure script of below

$ ./configure --enable-debug

Besides all the debugging advantages of building with this option, it will also be significantly speed up the build (at the cost of slower built programs). So when you are testing something small or working on the build system itself, it will be much faster to test your work with this option.


Build all make check tests within Valgrind. For more, see the description of --enable-check-with-valgrind in Gnuastro configure options.

-j INT
--jobs INT

The maximum number of threads/jobs for Make to build at any moment. As the name suggests (Make has an identical option), the number given to this option is directly passed on to any call of Make with its -j option.


After finishing the build, also run make check. By default, make check isn’t run because the developer usually has their own checks to work on (for example defined in tests/


After finishing the build, also run make install.


Run make dist-lzip pdf to build a distribution tarball (in .tar.lz format) and a PDF manual. This can be useful for archiving, or sending to colleagues who don’t use Git for an easy build and manual.

-u STR
--upload STR

Activate the --dist (-D) option, but also rename the tarball suffix to -latest.tar.lz (instead of the version number). Then use secure copy (scp, part of the SSH tools) to copy the tarball and PDF to the server and directory specified in the value to this option. For example --upload my-server:dir, will copy the two files to the dir directory of my-server server.


Short for --autoreconf --clean --debug --check --upload STR. --debug is added because it will greatly speed up the build. It will have no effect on the produced tarball. This is good when you have made a commit and are ready to publish it on your server (if nothing crashes). Recall that if any of the previous steps fail the script aborts.


Short for --autoreconf --clean --check --install --dist. This is useful when you actually want to install the commit you just made (if the build and checks succeed). It will also produce a distribution tarball and PDF manual for easy access to the installed tarball on your system at a later time.

Ideally, Gnuastro’s Git version history makes it easy for a prepared system to revert back to a different point in history. But Gnuastro also needs to bootstrap files and also your collaborators might (usually do!) find it too much of a burden to do the bootstrapping themselves. So it is convenient to have a tarball and PDF manual of the version you have installed (and are using in your research) handily available.


Print a description of this script along with all the options and their current values.

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3.3.3 Tests

After successfully building (compiling) the programs with the $ make command you can check the installation before installing. To run the tests, run

$ make check

For every program some tests are designed to check some possible operations. Running the command above will run those tests and give you a final report. If everything is OK and you have built all the programs, all the tests should pass. In case any of the tests fail, please have a look at Known issues and if that still doesn’t fix your problem, look that the ./tests/test-suite.log file to see if the source of the error is something particular to your system or more general. If you feel it is general, please contact us because it might be a bug. Note that the tests of some programs depend on the outputs of other program’s tests, so if you have not installed them they might be skipped or fail. Prior to releasing every distribution all these tests are checked. If you have a reasonably modern terminal, the outputs of the successful tests will be colored green and the failed ones will be colored red.

These scripts can also act as a good set of examples for you to see how the programs are run. All the tests are in the tests/ directory. The tests for each program are shell scripts (ending with .sh) in a sub-directory of this directory with the same name as the program. See Test scripts for more detailed information about these scripts in case you want to inspect them.

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3.3.4 A4 print book

The default print version of this book is provided in the letter paper size. If you would like to have the print version of this book on paper and you are living in a country which uses A4, then you can rebuild the book. The great thing about the GNU build system is that the book source code which is in Texinfo is also distributed with the program source code, enabling you to do such customization (hacking).

In order to change the paper size, you will need to have GNU Texinfo installed. Open doc/gnuastro.texi with any text editor. This is the source file that created this book. In the first few lines you will see this line:


In Texinfo, a line is commented with @c. Therefore, uncomment this line by deleting the first two characters such that it changes to:


Save the file and close it. You can now run

$ make pdf

and the new PDF book will be available in SRCdir/doc/gnuastro.pdf. By changing the pdf in $ make pdf to ps or dvi you can have the book in those formats. Note that you can do this for any book that is in Texinfo format, they might not have @afourpaper line, so you can add it close to the top of the Texinfo source file.

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3.3.5 Known issues

Depending on your operating system and the version of the compiler you are using, you might confront some known problems during the configuration ($ ./configure), compilation ($ make) and tests ($ make check). Here, their solutions are discussed.

If your problem was not listed above, please file a bug report (Report a bug).

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4 Common program behavior

All the programs in Gnuastro share a set of common behavior mainly to do with user interaction to facilitate their usage. The most basic is how you can configure each program to do what you want: define the input, change parameter/option values, or identify the output. All Gnuastro programs can also read your desired configuration from pre-defined or user-specified files so you don’t have to specify all the (sometimes numerous) parameters on the command-line each time you run a program. These files define the “default” program behavior in each directory, for each user, or on each system. In other cases, some programs can greatly benefit from the many threads available in modern CPUs, so here we’ll also discuss how you can get the most out of your hardware. Among some other issues, we will also discuss how you can get immediate and distraction-free (without taking your hands off the keyboard!) help, or access to this whole book, on the command-line.

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4.1 Command-line

All the programs in Gnuastro are customized through the standard GNU style command-line options. Thus, we’ll start by defining this general style that is very common in many command-line tools on Unix-like operating systems. Finally, the options that are common to all the programs in Gnuastro are discussed.

The command-line text that you type is passed onto the shell (or program managing the command-line) as a string of characters. See the “Invoking ProgramName” sections in this manual for some examples of commands with each program, for example Invoking Table. That string is then broken up into separate tokens or words by any metacharacters (like space, tab, |, > or ;) that might exist in the text. To learn more, please see the GNU Bash manual, for the complete list of meta-characters and other GNU Bash definitions (GNU Bash is the most common shell program). Its “Shell Operation” section has a short summary of the steps the shell takes before passing the commands to the program you called.

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4.1.1 Arguments and options

On the command-line, the first thing you usually enter is the name of the program you want to run. After that, you can specify two types of input: arguments and options. In the GNU-style, arguments are those tokens that are not preceded by any hyphens (-, see Arguments). Here is one example:

$ astcrop --center=53.162551,-27.789676 -w10/3600 --mode=wcs udf.fits

In this example, the argument is udf.fits. Arguments are most commonly the input file names containing your data. Options start with one or two hyphens, followed by an identifier for the option (the option’s name) and its value (see Options). Through options you tell the program how to interpret the data. In this example, we are running Crop to crop a region of width 10 arc-seconds centered at the given RA and Dec from the input Hubble Ultra-Deep Field (UDF) FITS image. So options come with an identifier (the option name which is separate from their value).

Arguments can be both mandatory and optional and unlike options they don’t have any identifiers (or help from you). Hence, their order might also matter (for example in cp which is used for copying one file to another location). The outputs of --usage and --help shows which arguments are optional and which are mandatory, see --usage. As their name suggests, options on the command-line can be considered to be optional and most of the time, you don’t have to worry about what order you specify them in. When the order does matter, or the option can be invoked multiple times, it is explicitly mentioned in the “Invoking ProgramName” section of each program.

In case your arguments or option values contain any of the shell’s meta-characters, you have to quote them. If there is only one such character, you can use a backslash (\) before it. If there are multiple, it might be easier to simply put your whole argument or option value inside of double quotes ("). In such cases, everything inside the double quotes will be seen as one token or word.

For example, let’s say you want to specify the header data unit (HDU) of your FITS file using a complex expression like ‘3; images(exposure > 100)’. If you simply add these after the --hdu (-h) option, the programs in Gnuastro will read the value to the HDU option as ‘3’ and run. Then, Bash will attempt to run a separate command ‘images(exposure > 100)’ and complain about a syntax error. This is because the semicolon (;) is an ‘end of command’ character in the shell. To solve this problem you can simply put double quotes around the whole string you want to pass to --hdu as seen below:

$ astcrop --hdu="3; images(exposure > 100)" FITSimage.fits

Alternatively you can put a ‘\’ before every meta-character in this string, but try doing that, and probably you will agree that the double quotes are much more easier, elegant and readable.

Next: , Previous: , Up: Arguments and options   [Contents][Index] Arguments

In Gnuastro, arguments are almost exclusively used as the input data file names. Please consult the first few paragraph of the “Invoking ProgramName” section for each program for a description of what it expects as input, how many arguments, or input data, it accepts, or in what order. Everything particular about how a program treats arguments, is explained under the “Invoking ProgramName” section for that program.

Generally, if there is a standard file name extension for a particular format, that filename extension is used to separate the kinds of arguments. The list below shows the data formats that are recognized in Gnuastro’s programs based on their file name endings. Any argument that doesn’t end with the specified extensions below is considered to be a text file (usually catalogs, see Tables). In some cases, a program can accept specific formats, for example ConvertType also accepts .jpg images.

Through out this book and in the command-line outputs, whenever we want to generalize all such astronomical data formats in a text place holder, we will use ASTRdata, we will assume that the extension is also part of this name. Any file ending with these names is directly passed on to CFITSIO to read. Therefore you don’t necessarily have to have these files on your computer, they can also be located on an FTP or HTTP server too, see the CFITSIO manual for more information.

CFITSIO has its own error reporting techniques, if your input file(s) cannot be opened, or read, those errors will be printed prior to the final error by Gnuastro.

Previous: , Up: Arguments and options   [Contents][Index] Options

Command-line options allow configuring the behavior of a program in all GNU/Linux applications for each particular execution on a particular input data. A single option can be called in two ways: long or short. All options in Gnuastro accept the long format which has two hyphens an can have many characters (for example --hdu). Short options only have one hyphen (-) followed by one character (for example -h). You can see some examples in the list of options in Common options or those for each program’s “Invoking ProgramName” section. Both formats are shown for those which support both. First the short is shown then the long.

Usually, the short options are for when you are writing on the command-line and want to save keystrokes and time. The long options are good for shell scripts, where you aren’t usually rushing. Long options provide a level of documentation, since they are more descriptive and less cryptic. Usually after a few months of not running a program, the short options will be forgotten and reading your previously written script will not be easy.

Some options need to be given a value if they are called and some don’t. You can think of the latter type of options as on/off options. These two types of options can be distinguished using the output of the --help and --usage options, which are common to all GNU software, see Getting help. In Gnuastro we use the following strings to specify when the option needs a value and what format that value should be in. More specific tests will be done in the program and if the values are out of range (for example negative when the program only wants a positive value), an error will be reported.


The value is read as an integer.


The value is read as a float. There are generally two types, depending on the context. If they are for fractions, they will have to be less than or equal to unity.


The value is read as a string of characters (for example a file name) or other particular settings like a HDU name, see below.

To specify a value in the short format, simply put the value after the option. Note that since the short options are only one character long, you don’t have to type anything between the option and its value. For the long option you either need white space or an = sign, for example -h2, -h 2, --hdu 2 or --hdu=2 are all equivalent.

The short format of on/off options (those that don’t need values) can be concatenated for example these two hypothetical sequences of options are equivalent: -a -b -c4 and -abc4. As an example, consider the following command to run Crop:

$ astcrop -Dr3 --wwidth 3 catalog.txt --deccol=4 ASTRdata

The $ is the shell prompt, astcrop is the program name. There are two arguments (catalog.txt and ASTRdata) and four options, two of them given in short format (-D, -r) and two in long format (--width and --deccol). Three of them require a value and one (-D) is an on/off option.

If an abbreviation is unique between all the options of a program, the long option names can be abbreviated. For example, instead of typing --printparams, typing --print or maybe even --pri will be enough, if there are conflicts, the program will warn you and show you the alternatives. Finally, if you want the argument parser to stop parsing arguments beyond a certain point, you can use two dashes: --. No text on the command-line beyond these two dashes will be parsed.

Gnuastro has two types of options with values, those that only take a single value are the most common type. If these options are repeated or called more than once on the command-line, the value of the last time it was called will be assigned to it. This is very useful when you are testing/experimenting. Let’s say you want to make a small modification to one option value. You can simply type the option with a new value in the end of the command and see how the script works. If you are satisfied with the change, you can remove the original option for human readability. If the change wasn’t satisfactory, you can remove the one you just added and not worry about forgetting the original value. Without this capability, you would have to memorize or save the original value somewhere else, run the command and then change the value again which is not at all convenient and is potentially cause lots of bugs.

On the other hand, some options can be called multiple times in one run of a program and can thus take multiple values (for example see the --column option in Invoking Table. In these cases, the order of stored values is the same order that you specified on the command-line.

Gnuastro’s programs don’t keep any internal default values, so some options are mandatory and if they don’t have a value, the program will complain and abort. Most programs have many such options and typing them by hand on every call is impractical. To facilitate the user experience, after parsing the command-line, Gnuastro’s programs read special configuration files to get the necessary values for the options you haven’t identified on the command-line. These configuration files are fully described in Configuration files.

CAUTION: In specifying a file address, if you want to use the shell’s tilde expansion (~) to specify your home directory, leave at least one space between the option name and your value. For example use -o ~/test, --output ~/test or --output= ~/test. Calling them with -o~/test or --output=~/test will disable shell expansion.

CAUTION: If you forget to specify a value for an option which requires one, and that option is the last one, Gnuastro will warn you. But if it is in the middle of the command, it will take the text of the next option or argument as the value which can cause undefined behavior.

NOTE: In some contexts Gnuastro’s counting starts from 0 and in others 1. You can assume by default that counting starts from 1, if it starts from 0 for a special option, it will be explicitly mentioned.

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4.1.2 Common options

To facilitate the job of the users and developers, all the programs in Gnuastro share some basic command-line options for the options that are common to many of the programs. The full list is classified as Input/Output options, Processing options, and Operating mode options. In some programs, some of the options are irrelevant, but still recognized (you won’t get an unrecognized option error, but the value isn’t used). Unless otherwise mentioned, these options are identical between all programs.

Next: , Previous: , Up: Common options   [Contents][Index] Input/Output options

These options are to do with the input and outputs of the various programs.


The name or number of the desired Header Data Unit, or HDU, in the FITS image. A FITS file can store multiple HDUs or extensions, each with either an image or a table or nothing at all (only a header). Note that counting of the extensions starts from 0(zero), not 1(one). Counting from 0 is forced on us by CFITSIO which directly reads the value you give with this option (see CFITSIO). When specifying the name, case is not important so IMAGE, image or ImAgE are equivalent.

CFITSIO has many capabilities to help you find the extension you want, far beyond the simple extension number and name. See CFITSIO manual’s “HDU Location Specification” section for a very complete explanation with several examples. A # is appended to the string you specify for the HDU74 and the result is put in square brackets and appended to the FITS file name before calling CFITSIO to read the contents of the HDU for all the programs in Gnuastro.

-s STR

Where to match/search for columns when the column identifier wasn’t a number, see Selecting table columns. The acceptable values are name, unit, or comment. This option is only relevant for programs that take table columns as input.


Ignore case while matching/searching column meta-data (in the field specified by the --searchin). The FITS standard suggests to treat the column names as case insensitive, which is strongly recommended here also but is not enforced. This option is only relevant for programs that take table columns as input.

This option is not relevant to BuildProgram, hence in that program the short option -I is used for include directories, not to ignore case.

-o STR

The name of the output file or directory. With this option the automatic output names explained in Automatic output are ignored.


The data type of the output depending on the program context. This option isn’t applicable to some programs like Fits and will be ignored by them. The different acceptable values to this option are fully described in Numeric data types.


By default, if the output file already exists, Gnuastro’s programs will silently delete it and put their own outputs in its place. When this option is activated, if the output file already exists, the programs will not delete it, will warn you, and will abort.


In automatic output names, don’t remove the directory information of the input file names. As explained in Automatic output, if no output name is specified (with --output), then the output name will be made in the existing directory based on your input’s file name (ignoring the directory of the input). If you call this option, the directory information of the input will be kept and the automatically generated output name will be in the same directory as the input (usually with a suffix added). Note that his is only relevant if you are running the program in a different directory than the input data.

-t STR

The output table’s type. This option is only relevant when the output is a table and its format cannot be deduced from its filename. For example, if a name ending in .fits was given to --output, then the program knows you want a FITS table. But there are two types of FITS tables: FITS ASCII, and FITS binary. Thus, with this option, the program is able to identify which type you want. The currently recognized values to this option are:


A plain text table with white-space characters between the columns (see Gnuastro text table format).


A FITS ASCII table (see Recognized table formats).


A FITS binary table (see Recognized table formats).

Next: , Previous: , Up: Common options   [Contents][Index] Processing options

Some processing steps are common to several programs, so they are defined as common options to all programs. Note that this class of common options is thus necessarily less common between all the programs than those described in Input/Output options, or Operating mode options options. Also, if they are irrelevant for a program, these options will not display in the --help output of the program.


The minimum size (in bytes) to store the contents of each main processing array of a program as a file (on the non-volatile HDD/SSD), not in RAM. This can be very useful for large datasets which can be very memory intensive such that your RAM will not be sufficient to keep/process them. A random filename is assigned to the array (in a .gnuastro directory within the running directory) which will keep the contents of the array as long as it is necessary.

When this option has a value of 0 (zero), all arrays that use this option in a program will actually be in a file (not in RAM). When the value is -1 (largest possible number in the unsigned integer types) these arrays will be definitely allocated in RAM. However, for complex programs like NoiseChisel, it is recommended to not set it to 0, but a value like 10000 so the many small arrays necessary during processing are stored in RAM and only larger ones are saved as a file.

Please note that using a non-volatile file (in the HDD/SDD) instead of RAM can significantly increase the program’s running time, especially on HDDs. So it is best to give this option large values by default. You can then decrease it for a specific program’s invocation on a large input after you see memory issues arise (for example an error, or the program not aborting and fully consuming your memory).

The random file will be deleted once it is no longer needed by the program. The .gnuastro directory will also be deleted if it has no other contents (you may also have configuration files in this directory, see Configuration files). If you see randomly named files remaining in this directory, please send us a bug report so we address the problem, see Report a bug.

-Z INT[,INT[,...]]

The size of regular tiles for tessellation, see Tessellation. For each dimension an integer length (in units of data-elements or pixels) is necessary. If the number of input dimensions is different from the number of values given to this option, the program will stop with an error. Values must be separated by commas (,) and can also be fractions (for example 4/2). If they are fractions, the result must be an integer, otherwise an error will be printed.

-M INT[,INT[,...]]

The number of channels for larger input tessellation, see Tessellation. The number and types of acceptable values are similar to --tilesize. The only difference is that instead of length, the integers values given to this option represent the number of channels, not their size.


The fraction of remainder size along all dimensions to add to the first tile. See Tessellation for a complete description. This option is only relevant if --tilesize is not exactly divisible by the input dataset’s size in a dimension. If the remainder size is larger than this fraction (compared to --tilesize), then the remainder size will be added with one regular tile size and divided between two tiles at the start and end of the given dimension.


Ignore the channel borders for the high-level job of the given application. As a result, while the channel borders are respected in defining the small tiles (such that no tile will cross a channel border), the higher-level program operation will ignore them, see Tessellation.


Make a FITS file with the same dimensions as the input but each pixel is replaced with the ID of the tile that it is associated with. Note that the tile IDs start from 0. See Tessellation for more on Tiling an image in Gnuastro.


When showing the tile values (for example with --checktiles, or when the program’s output is tessellated) only use one element for each tile. This can be useful when only the relative values given to each tile compared to the rest are important or need to be checked. Since the tiles usually have a large number of pixels within them the output will be much smaller, and so easier to read, write, store, or send.

Note that when the full input size in any dimension is not exactly divisible by the given --tilesize in that dimension, the edge tile(s) will have different sizes (in units of the input’s size), see --remainderfrac. But with this option, all displayed values are going to have the (same) size of one data-element. Hence, in such cases, the image proportions are going to be slightly different with this option.

If your input image is not exactly divisible by the tile size and you want one value per tile for some higher-level processing, all is not lost though. You can see how many pixels were within each tile (for example to weight the values or discard some for later processing) with Gnuastro’s Statistics (see Statistics) as shown below. The output FITS file is going to have two extensions, one with the median calculated on each tile and one with the number of elements that each tile covers. You can then use the where operator in Arithmetic to set the values of all tiles that don’t have the regular area to a blank value.

$ aststatistics --median --number --ontile input.fits    \
                --oneelempertile --output=o.fits
$ REGULAR_AREA=1600    # Check second extension of `o.fits'.
$ astarithmetic o.fits o.fits $REGULAR_AREA ne nan where \
                -h1 -h2

Note that if input.fits also has blank values, then the median on tiles with blank values will also be ignored with the command above (which is desirable).


When values are to be interpolated, only change the values of the blank elements, keep the non-blank elements untouched.


The number of nearby non-blank neighbors to use for interpolation.

Previous: , Up: Common options   [Contents][Index] Operating mode options

Another group of options that are common to all the programs in Gnuastro are those to do with the general operation of the programs. The explanation for those that are not only limited to Gnuastro but are common to all GNU programs start with (GNU option).


(GNU option) Stop parsing the command-line. This option can be useful in scripts or when using the shell history. Suppose you have a long list of options, and want to see if removing some of them (to read from configuration files, see Configuration files) can give a better result. If the ones you want to remove are the last ones on the command-line, you don’t have to delete them, you can just add -- before them and if you don’t get what you want, you can remove the -- and get the same initial result.


(GNU option) Only print the options and arguments and abort. This is very useful for when you know the what the options do, and have just forgot their long/short identifiers, see --usage.


(GNU option) Print all options with an explanation and abort. Adding this option will print all the options in their short and long formats, also displaying which ones need a value if they are called (with an = after the long format followed by a string specifying the format, see Options). A short explanation is also given for what the option is for. The program will quit immediately after the message is printed and will not do any form of processing, see --help.


(GNU option) Print a short message, showing the full name, version, copyright information and program authors and abort. On the first line, it will print the official name (not executable name) and version number of the program. Following this is a blank line and a copyright information. The program will not run.


Don’t report steps. All the programs in Gnuastro that have multiple major steps will report their steps for you to follow while they are operating. If you do not want to see these reports, you can call this option and only error/warning messages will be printed. If the steps are done very fast (depending on the properties of your input) disabling these reports will also decrease running time.


Print the BibTeX entry for Gnuastro and the particular program (if that program comes with a separate paper) and abort. Citations are vital for the continued work on Gnuastro. Gnuastro started and is continued based on separate research projects. So if you find any of the tools offered in Gnuastro to be useful in your research, please use the output of this command to cite the program and Gnuastro in your research paper. Thank you.

Gnuastro is still new, there is no separate paper only devoted to Gnuastro yet. Therefore currently the paper to cite for Gnuastro is the paper for NoiseChisel which is the first published paper introducing Gnuastro to the astronomical community. Upon reaching a certain point, a paper completely devoted to Gnuastro will be published, see GNU Astronomy Utilities 1.0.


With this option, Gnuastro’s programs will read your command-line options and all the configuration files. If there is no problem (like a missing parameter or a value in the wrong format or range) and immediately before actually running, the programs will print the full list of option names, values and descriptions, sorted and grouped by context and abort. They will also report the version number, the date they were configured on your system and the time they were reported.

As an example, you can give your full command-line options and even the input and output file names and finally just add -P to check if all the parameters are finely set. If everything is OK, you can just run the same command (easily retrieved from the shell history, with the top arrow key) and simply remove the last two characters that showed this option.

Since no program will actually start its processing when this option is called, the otherwise mandatory arguments for each program (for example input image or catalog files) are no longer required when you call this option.


Parse STR as a configuration file immediately when this option is confronted (see Configuration files). The --config option can be called multiple times in one run of any Gnuastro program on the command-line or in the configuration files. In any case, it will be immediately read (before parsing the rest of the options on the command-line, or lines in a configuration file).

Note that by definition, later options on the command-line still take precedence over those in these in any configuration file, including the file(s) given to this option. Also see --lastconfig and --onlyversion on how this option can be used for reproducible results.


Update the current directory configuration file for the Gnuastro program and quit. The full set of command-line and configuration file options will be parsed and options with a value will be written in the current directory configuration file for this program (see Configuration files). If the configuration file or its directory doesn’t exist, it will be created. If a configuration file exists it will be replaced (after it, and all other configuration files have been read). In any case, the program will not run.

This is the recommended method75 to edit/set the configuration file for all future calls to Gnuastro’s programs. It will internally check if your values are in the correct range and type and save them according to the configuration file format, see Configuration file format. So if there are unreasonable values to some options, the program will notify you and abort before writing the final configuration file.

When this option is called, the otherwise mandatory arguments, for example input image or catalog file(s), are no longer mandatory (since the program will not run).


Update the user configuration file and quit (see Configuration files). See explanation under --setdirconf for more details.


This is the last configuration file that must be read. When this option is confronted in any stage of reading the options (on the command-line or in a configuration file), no other configuration file will be parsed, see Configuration file precedence and Current directory and User wide. Like all on/off options, on the command-line, this option doesn’t take any values. But in a configuration file, it takes the values of 0 or 1, see Configuration file format. If it is present in a configuration file with a value of 0, then all later occurrences of this option will be ignored.


Only run the program if Gnuastro’s version is exactly equal to STR (see Version numbering). Note that it is not compared as a number, but as a string of characters, so 0, or 0.0 and 0.00 are different. If the running Gnuastro version is different, then this option will report an error and abort as soon as it is confronted on the command-line or in a configuration file. If the running Gnuastro version is the same as STR, then the program will run as if this option was not called.

This is useful if you want your results to be exactly reproducible and not mistakenly run with an updated/newer or older version of the program. Besides internal algorithmic/behavior changes in programs, the existence of options or their names might change between versions (especially in these earlier versions of Gnuastro).

Hence, when using this option (probably in a script or in a configuration file), be sure to call it before other options. The benefit is that, when the version differs, the other options won’t be parsed and you, or your collaborators/users, won’t get errors saying an option in your configuration doesn’t exist in the running version of the program.

Here is one example of how this option can be used in conjunction with the --lastconfig option. Let’s assume that you were satisfied with the results of this command: astnoisechisel image.fits --snquant=0.95 (along with various options set in various configuration files). You can save the state of NoiseChisel and reproduce that exact result on image.fits later by following these steps (the the extra spaces, and \, are only for easy readability, if you want to try it out, only one space between each token is enough).

$ echo "onlyversion X.XX"             > reproducible.conf
$ echo "lastconfig 1"                >> reproducible.conf
$ astnoisechisel image.fits --snquant=0.95 -P            \
                                     >> reproducible.conf

--onlyversion was available from Gnuastro 0.0, so putting it immediately at the start of a configuration file will ensure that later, you (or others using different version) won’t get a non-recognized option error in case an option was added/removed. --lastconfig will inform the installed NoiseChisel to not parse any other configuration files. This is done because we don’t want the user’s user-wide or system wide option values affecting our results. Finally, with the third command, which has a -P (short for --printparams), NoiseChisel will print all the option values visible to it (in all the configuration files) and the shell will append them to reproduce.conf. Hence, you don’t have to worry about remembering the (possibly) different options in the different configuration files.

Afterwards, if you run NoiseChisel as shown below (telling it to read this configuration file with the --config option). You can be sure that there will either be an error (for version mismatch) or it will produce exactly the same result that you got before.

$ astnoisechisel --config=reproducible.conf

Some programs can generate extra information about their outputs in a log file. When this option is called in those programs, the log file will also be printed. If the program doesn’t generate a log file, this option is ignored.

--log isn’t thread-safe: The log file usually has a fixed name. Therefore if two simultaneous calls (with --log) of a program are made in the same directory, the program will try to write to the same file. This will cause problems like unreasonable log file, undefined behavior, or a crash.


Use INT CPU threads when running a Gnuastro program (see Multi-threaded operations). If the value is zero (0), or this option is not given on the command-line or any configuration file, the value will be determined at run-time: the maximum number of threads available to the system when you run a Gnuastro program.

Note that multi-threaded programming is only relevant to some programs. In others, this option will be ignored.

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4.2 Configuration files

Each program needs a certain number of parameters to run. Supplying all the necessary parameters each time you run the program is very frustrating and prone to errors. Therefore all the programs read the values for the necessary options you have not given in the command line from one of several plain text files (which you can view and edit with any text editor). These files are known as configuration files and are usually kept in a directory named etc/ according to the file system hierarchy standard76.

The thing to have in mind is that none of the programs in Gnuastro keep any internal default value. All the values must either be stored in one of the configuration files or explicitly called in the command-line. In case the necessary parameters are not given through any of these methods, the program will print a missing option error and abort. The only exception to this is --numthreads, whose default value is determined at run-time using the number of threads available to your system, see Multi-threaded operations. Of course, you can still provide a default value for the number of threads at any of the levels below, but if you don’t, the program will not abort. Also note that through automatic output name generation, the value to the --output option is also not mandatory on the command-line or in the configuration files for all programs which don’t rely on that value as an input77, see Automatic output.

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4.2.1 Configuration file format

The configuration files for each program have the standard program executable name with a ‘.conf’ suffix. When you download the source code, you can find them in the same directory as the source code of each program, see Program source.

Any line in the configuration file whose first non-white character is a # is considered to be a comment and is ignored. An empty line is also similarly ignored. The long name of the option should be used as an identifier. The parameter name and parameter value have to be separated by any number of ‘white-space’ characters: space, tab or vertical tab. By default several space characters are used. If the value of an option has space characters (most commonly for the hdu option), then the full value can be enclosed in double quotation signs (", similar to the example in Arguments and options). If it is an option without a value in the --help output (on/off option, see Options), then the value should be 1 if it is to be ‘on’ and 0 otherwise.

In each non-commented and non-blank line, any text after the first two words (option identifier and value) is ignored. If an option identifier is not recognized in the configuration file, the name of the file, the line number of the unrecognized option, and the unrecognized identifier name will be reported and the program will abort. If a parameter is repeated more more than once in the configuration files, accepts only one value, and is not set on the command-line, then only the first value will be used, the rest will be ignored.

You can build or edit any of the directories and the configuration files yourself using any text editor. However, it is recommended to use the --setdirconf and --setusrconf options to set default values for the current directory or this user, see Operating mode options. With these options, the values you give will be checked before writing in the configuration file. They will also print a set of commented lines guiding the reader and will also classify the options based on their context and write them in their logical order to be more understandable.

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4.2.2 Configuration file precedence

The option values in all the programs of Gnuastro will be filled in the following order. If an option only takes one value which is given in an earlier step, any value for that option in a later step will be ignored. Note that if the lastconfig option is specified in any step below, all later files will be ignored (see Operating mode options). The basic idea behind setting this progressive state of checking for parameter values is that separate users of a computer or separate folders in a user’s file system might need different values for some parameters.

In each step, there can also be a configuration file containing the common options in all the programs: gnuastro.conf (see Common options). If options specific to one program are specified in this file, there will be unrecognized option errors, or unexpected behavior if the option has different behavior in another program. On the other hand, there is no problem with astprogname.conf containing common options78.

  1. Command-line options, for a particular run of ProgramName.
  2. .gnuastro/astprogname.conf is parsed by ProgramName in the current directory.
  3. .gnuastro/gnuastro.conf is parsed by all Gnuastro programs in the current directory.
  4. $HOME/.local/etc/astprogname.conf is parsed by ProgramName in the user’s home directory (see Current directory and User wide).
  5. $HOME/.local/etc/gnuastro.conf is parsed by all Gnuastro programs in the user’s home directory (see Current directory and User wide).
  6. prefix/etc/astprogname.conf is parsed by ProgramName in the system-wide installation directory (see System wide for prefix).
  7. prefix/etc/gnuastro.conf is parsed by all Gnuastro programs in the system-wide installation directory (see System wide for prefix).

Manipulating the order: You can manipulate this order or add new files with the following two options which are fully described in Operating mode options:


Allows you to define any file to be parsed as a configuration file on the command-line or within the any other configuration file. Recall that the file given to --config is parsed immediately when this option is confronted (on the command-line or in a configuration file).


Allows you to stop the parsing of subsequent configuration files. Note that if this option is given in a configuration file, it will be fully read, so its position in the configuration doesn’t matter (unlike --config).

One example of benefiting from these configuration files can be this: raw telescope images usually have their main image extension in the second FITS extension, while processed FITS images usually only have one extension. If your system-wide default input extension is 0 (the first), then when you want to work with the former group of data you have to explicitly mention it to the programs every time. With this progressive state of default values to check, you can set different default values for the different directories that you would like to run Gnuastro in for your different purposes, so you won’t have to worry about this issue any more.

The same can be said about the gnuastro.conf files: by specifying a behavior in this single file, all Gnuastro programs in the respective directory, user, or system-wide steps will behave similarly. For example to keep the input’s directory when no specific output is given (see Automatic output), or to not delete an existing file if it has the same name as a given output (see Input/Output options).

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4.2.3 Current directory and User wide

For the current (local) and user-wide directories, the configuration files are stored in the hidden sub-directories named .gnuastro/ and $HOME/.local/etc/ respectively. Unless you have changed it, the $HOME environment variable should point to your home directory. You can check it by running $ echo $HOME. Each time you run any of the programs in Gnuastro, this environment variable is read and placed in the above address. So if you suddenly see that your home configuration files are not being read, probably you (or some other program) has changed the value of this environment variable.

Although it might cause confusions like above, this dependence on the HOME environment variable enables you to temporarily use a different directory as your home directory. This can come in handy in complicated situations. To set the user or current directory configuration files based on your command-line input, you can use the --setdirconf or --setusrconf, see Operating mode options.

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4.2.4 System wide

When Gnuastro is installed, the configuration files that are shipped with the distribution are copied into the (possibly system wide) prefix/etc/ directory. For more details on prefix, see Installation directory (by default it is: /usr/local). This directory is the final place (with the lowest priority) that the programs in Gnuastro will check to retrieve parameter values.

If you remove an option and its value from the system wide configuration files, you either have to specify it in more immediate configuration files or set it each time in the command-line. Recall that none of the programs in Gnuastro keep any internal default values and will abort if they don’t find a value for the necessary parameters (except the number of threads and output file name). So even though you might never expect to use an optional option, it safe to have it available in this system-wide configuration file even if you don’t intend to use it frequently.

Note that in case you install Gnuastro from your distribution’s repositories, prefix will either be set to / (the root directory) or /usr, so you can find the system wide configuration variables in /etc/ or /usr/etc/. The prefix of /usr/local/ is conventionally used for programs you install from source by your self as in Quick start.

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4.3 Multi-threaded operations

Some of the programs benefit significantly when you use all the threads your computer’s CPU has to offer to your operating system. The number of threads available can be larger than the number of physical (hardware) cores in the CPU (also known as Simultaneous multithreading). For example, in Intel’s CPUs (those that implement its Hyper-threading technology) the number of threads is usually double the number of physical cores in your CPU. On a GNU/Linux system, the number of threads available can be found with the command $ nproc command (part of GNU Coreutils).

Gnuastro’s programs can find the number of threads available to your system internally at run-time (when you execute the program). However, if a value is given to the --numthreads option, the given number will be used, see Operating mode options and Configuration files for ways to use this option. Thus --numthreads is the only common option in Gnuastro’s programs with a value that doesn’t have to be specified anywhere on the command-line or in the configuration files.

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4.3.1 A note on threads

Spinning off threads is not necessarily the most efficient way to run an application. Creating a new thread isn’t a cheap operation for the operating system. It is most useful when the input data are fixed and you want the same operation to be done on parts of it. For example one input image to Crop and multiple crops from various parts of it. In this fashion, the image is loaded into memory once, all the crops are divided between the number of threads internally and each thread cuts out those parts which are assigned to it from the same image. On the other hand, if you have multiple images and you want to crop the same region(s) out of all of them, it is much more efficient to set --numthreads=1 (so no threads spin off) and run Crop multiple times simultaneously, see How to run simultaneous operations.

You can check the boost in speed by first running a program on one of the data sets with the maximum number of threads and another time (with everything else the same) and only using one thread. You will notice that the wall-clock time (reported by most programs at their end) in the former is longer than the latter divided by number of physical CPU cores (not threads) available to your operating system. Asymptotically these two times can be equal (most of the time they aren’t). So limiting the programs to use only one thread and running them independently on the number of available threads will be more efficient.

Note that the operating system keeps a cache of recently processed data, so usually, the second time you process an identical data set (independent of the number of threads used), you will get faster results. In order to make an unbiased comparison, you have to first clean the system’s cache with the following command between the two runs.

$ sync; echo 3 | sudo tee /proc/sys/vm/drop_caches

SUMMARY: Should I use multiple threads? Depends:

  • If you only have one data set (image in most cases!), then yes, the more threads you use (with a maximum of the number of threads available to your OS) the faster you will get your results.
  • If you want to run the same operation on multiple data sets, it is best to set the number of threads to 1 and use Make, or GNU Parallel, as explained in How to run simultaneous operations.

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4.3.2 How to run simultaneous operations

There are two79 approaches to simultaneously execute a program: using GNU Parallel or Make (GNU Make is the most common implementation). The first is very useful when you only want to do one job multiple times and want to get back to your work without actually keeping the command you ran. The second is usually for more important operations, with lots of dependencies between the different products (for example a full scientific research).

GNU Parallel

When you only want to run multiple instances of a command on different threads and get on with the rest of your work, the best method is to use GNU parallel. Surprisingly GNU Parallel is one of the few GNU packages that has no Info documentation but only a Man page, see Info. So to see the documentation after installing it please run

$ man parallel

As an example, let’s assume we want to crop a region fixed on the pixels (500, 600) with the default width from all the FITS images in the ./data directory ending with sci.fits to the current directory. To do this, you can run:

$ parallel astcrop --numthreads=1 --xc=500 --yc=600 ::: \

GNU Parallel can help in many more conditions, this is one of the simplest, see the man page for lots of other examples. For absolute beginners: the backslash (\) is only a line breaker to fit nicely in the page. If you type the whole command in one line, you should remove it.


Make is a program for building “targets” (e.g., files) using “recipes” (a set of operations) when their known “prerequisites” (other files) have been updated. It elegantly allows you to define dependency structures for building your final output and updating it efficiently when the inputs change. It is the most common infra-structure to build software today.

Scientific research methodology is very similar to software development: you start by testing a hypothesis on a small sample of objects/targets with a simple set of steps. As you are able to get promising results, you improve the method and use it on a larger, more general, sample. In the process, you will confront many issues that have to be corrected (bugs in software development jargon). Make a wonderful tool to manage this style of development. It has been used to make reproducible papers, for example see the reproduction pipeline of the paper introducing NoiseChisel (one of Gnuastro’s programs).

GNU Make80 is the most common implementation which (similar to nearly all GNU programs, comes with a wonderful manual81). Make is very basic and simple, and thus the manual is short (the most important parts are in the first roughly 100 pages) and easy to read/understand.

Make comes with a --jobs (-j) option which allows you to specify the maximum number of jobs that can be done simultaneously. For example if you have 8 threads available to your operating system. You can run:

$ make -j8

With this command, Make will process your Makefile and create all the targets (can be thousands of FITS images for example) simultaneously on 8 threads, while fully respecting their dependencies (only building a file/target when its prerequisites are successfully built). Make is thus strongly recommended for managing scientific research where robustness, archiving, reproducibility and speed82 are important.

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4.4 Numeric data types

At the lowest level, the computer stores everything in terms of 1 or 0. For example, each program in Gnuastro, or each astronomical image you take with the telescope is actually a string of millions of these zeros and ones. The space required to keep a zero or one is the smallest unit of storage, and is known as a bit. However, understanding and manipulating this string of bits is extremely hard for most people. Therefore, we define packages of these bits along with a standard on how to interpret the bits in each package as a type.

The most basic standard for reading the bits is integer numbers (\(..., -2, -1, 0, 1, 2, ...\), more bits will give larger limits). The common integer types are 8, 16, 32, and 64 bits wide. For each width, there are two standards for reading the bits: signed and unsigned integers. In the former, negative numbers are allowed and in the latter, they aren’t. The unsigned types thus have larger positive limits (one extra bit), but no negative value. When the context of your work doesn’t involve negative numbers (for example counting, where negative is not defined), it is best to use the unsigned types. For full numerical range of all integer types, see below.

Another standard of converting a given number of bits to numbers is the floating point standard, this standard can approximately store any real number with a given precision. There are two common floating point types: 32-bit and 64-bit, for single and double precision floating point numbers respectively. The former is sufficient for data with less than 8 significant decimal digits (most astronomical data), while the latter is good for less than 16 significant decimal digits. The representation of real numbers as bits is much more complex than integers. If you are interested, you can start with the Wikipedia article.

With the conversion operators in Gnuastro’s Arithmetic, you can change the types of data to each other, which is necessary in some contexts. For example the program/library, that you intend to feed the data into, only accepts floating point values, but you have an integer image. Another situation that conversion can be helpful is when you know that your data only has values that fit within int8 or uint16. However it is currently formatted in the float64 type. Operations involving floating point or larger integer types are significantly slower than integer or smaller-width types respectively. In the latter case, it also requires much more (by 8 or 4 times in the example above) storage space. So when you confront such situations and want to store/archive/transfter the data, it is best convert them to the most efficient type.

The short and long names for the recognized numeric data types in Gnuastro are listed below. Both short and long names can be used when you want to specify a type. For example, as a value to the common option --type (see Input/Output options), or in the information comment lines of Gnuastro text table format. The ranges listed below are inclusive.


8-bit unsigned integers, range:
\([0\rm{\ to\ }2^8-1]\) or \([0\rm{\ to\ }255]\).


8-bit signed integers, range:
\([-2^7\rm{\ to\ }2^7-1]\) or \([-128\rm{\ to\ }127]\).


16-bit unsigned integers, range:
\([0\rm{\ to\ }2^{16}-1]\) or \([0\rm{\ to\ }65535]\).


16-bit signed integers, range:
\([-2^{15}\rm{\ to\ }2^{15}-1]\) or \([-32768\rm{\ to\ }32767]\).


32-bit unsigned integers, range:
\([0\rm{\ to\ }2^{32}-1]\) or \([0\rm{\ to\ }4294967295]\).


32-bit signed integers, range:
\([-2^{31}\rm{\ to\ }2^{31}-1]\) or \([-2147483648\rm{\ to\ }2147483647]\).


64-bit unsigned integers, range
\([0\rm{\ to\ }2^{64}-1]\) or \([0\rm{\ to\ }18446744073709551615]\).


64-bit signed integers, range:
\([-2^{63}\rm{\ to\ }2^{63}-1]\) or \([-9223372036854775808\rm{\ to\ }9223372036854775807]\).


32-bit (single-precision) floating point types. The maximum (minimum is its negative) possible value is \(3.402823\times10^{38}\). Single-precision floating points can accurately represent a floating point number up to \(\sim7.2\) significant decimals. Given the heavy noise in astronomical data, this is usually more than sufficient for storing results.


64-bit (double-precision) floating point types. The maximum (minimum is its negative) possible value is \(\sim10^{308}\). Double-precision floating points can accurately represent a floating point number \(\sim15.9\) significant decimals. This is usually good for processing (mixing) the data internally, for example a sum of single precision data (and later storing the result as float32).

Some file formats don’t recognize all types. Some file formats don’t recognize all the types, for example the FITS standard (see Fits) does not define uint64 in binary tables or images. When a type is not acceptable for output into a given file format, the respective Gnuastro program or library will let you know and abort. On the command-line, you can use the Arithmetic program to convert the numerical type of a dataset, in the libraries, you can call gal_data_copy_to_new_type.

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4.5 Tables

“A table is a collection of related data held in a structured format within a database. It consists of columns, and rows.” (from Wikipedia). Each column in the table contains the values of one property and each row is a collection of properties (columns) for one target object. For example, let’s assume you have just ran MakeCatalog (see MakeCatalog) on an image to measure some properties for the labeled regions (which might be detected galaxies for example) in the image. For each labeled region (detected galaxy), there will be a row which groups its measured properties as columns, one column for each property. One such property can be the object’s magnitude, which is the sum of pixels with that label, or its center can be defined as the light-weighted average value of those pixels. Many such properties can be derived from the raw pixel values and their position, see Invoking MakeCatalog for a long list.

As a summary, for each labeled region (or, galaxy) we have one row and for each measured property we have one column. This high-level structure is usually the first step for higher-level analysis, for example finding the stellar mass or photometric redshift from magnitudes in multiple colors. Thus, tables are not just outputs of programs, in fact it is much more common for tables to be inputs of programs. For example, to make a mock galaxy image, you need to feed in the properties of each galaxy into MakeProfiles for it do the inverse of the process above and make a simulated image from a catalog, see Sufi simulates a detection. In other cases, you can feed a table into Crop and it will crop out regions centered on the positions within the table, see Hubble visually checks and classifies his catalog. So to end this relatively long introduction, tables play a very important role in astronomy, or generally all branches of data analysis.

In Recognized table formats the currently recognized table formats in Gnuastro are discussed. You can use any of these tables as input or ask for them to be built as output. The most common type of table format is a simple plain text file with each row on one line and columns separated by white space characters, this format is easy to read/write by eye/hand. To give it the full functionality of more specific table types like the FITS tables, Gnuastro has a special convention which you can use to give each column a name, type, unit, and comments, while still being readable by other plain text table readers. This convention is described in Gnuastro text table format.

When tables are input to a program, the program reading it needs to know which column(s) it should use for its desired purposes. Gnuastro’s programs all follow a similar convention, on the way you can select columns in a table. They are thoroughly discussed in Selecting table columns.

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4.5.1 Recognized table formats

The list of table formats that Gnuastro can currently read from and write to are described below. Each has their own advantage and disadvantages, so a short review of the format is also provided to help you make the best choice based on how you want to define your input tables or later use your output tables.

Plain text table

This is the most basic and simplest way to create, view, or edit the table by hand on a text editor. The other formats described below are less eye-friendly and have a more formal structure (for easier computer readability). It is fully described in Gnuastro text table format.


The FITS ASCII table extension is fully in ASCII encoding and thus easily readable on any text editor (assuming it is the only extension in the FITS file). If the FITS file also contains binary extensions (for example an image or binary table extensions), then there will be many hard to print characters. The FITS ASCII format doesn’t have new line characters to separate rows. In the FITS ASCII table standard, each row is defined as a fixed number of characters (value to the NAXIS1 keyword), so to visually inspect it properly, you would have to adjust your text editor’s width to this value. All columns start at given character positions and have a fixed width (number of characters).

Numbers in a FITS ASCII table are printed into ASCII format, they are not in binary (that the CPU uses). Hence, they can take a larger space in memory, loose their precision, and take longer to read into memory. If you are dealing with integer type columns (see Numeric data types), another issue with FITS ASCII tables is that the type information for the column will be lost (there is only one integer type in FITS ASCII tables). One problem with the binary format on the other hand is that it isn’t portable (different CPUs/compilers) have different standards for translating the zeros and ones. But since ASCII characters are defined on a byte and are well recognized, they are better for portability on those various systems. Gnuastro’s plain text table format described below is much more portable and easier to read/write/interpret by humans manually.

Generally, as the name implies, this format is useful for when your table mainly contains ASCII columns (for example file names, or descriptions). They can be useful when you need to include columns with structured ASCII information along with other extensions in one FITS file. In such cases, you can also consider header keywords (see Fits).

FITS binary tables

The FITS binary table is the FITS standard’s solution to the issues discussed with keeping numbers in ASCII format as described under the FITS ASCII table title above. Only columns defined as a string type (a string of ASCII characters) are readable in a text editor. The portability problem with binary formats discussed above is mostly solved thanks to the portability of CFITSIO (see CFITSIO) and the very long history of the FITS format which has been widely used since the 1970s.

In the case of most numbers, storing them in binary format is more memory efficient than ASCII format. For example, to store -25.72034 in ASCII format, you need 9 bytes/characters. But if you keep this same number (to the approximate precision possible) as a 4-byte (32-bit) floating point number, you can keep/transmit it with less than half the amount of memory. When catalogs contain thousands/millions of rows in tens/hundreds of columns, this can lead to significant improvements in memory/band-width usage. Moreover, since the CPU does its operations in the binary formats, reading the table in and writing it out is also much faster than an ASCII table.

When you are dealing with integer numbers, the compression ratio can be even better, for example if you know all of the values in a column are positive and less than 255, you can use the unsigned char type which only takes one byte! If they are between -128 and 127, then you can use the (signed) char type. So if you are thoughtful about the limits of your integer columns, you can greatly reduce the size of your file and also the speed at which it is read/written. This can be very useful when sharing your results with collaborators or publishing them. To decrease the file size even more you can name your output as ending in .fits.gz so it is also compressed after creation. Just note that compression/decompressing is CPU intensive and can slow down the writing/reading of the file.

Fortunately the FITS Binary table format also accepts ASCII strings as column types (along with the various numerical types). So your dataset can also contain non-numerical columns.

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4.5.2 Gnuastro text table format

Plain text files are the most generic, portable, and easiest way to (manually) create, (visually) inspect, or (manually) edit a table. In this format, the ending of a row is defined by the new-line character (a line on a text editor). So when you view it on a text editor, every row will occupy one line. The delimiters (or characters separating the columns) are white space characters (space, horizontal tab, vertical tab) and a comma (,). The only further requirement is that all rows/lines must have the same number of columns.

The columns don’t have to be exactly under each other and the rows can be arbitrarily long with different lengths. For example the following contents in a file would be interpreted as a table with 4 columns and 2 rows, with each element interpreted as a double type (see Numeric data types).

1     2.234948   128   39.8923e8
2 , 4.454        792     72.98348e7

However, the example above has no other information about the columns (it is just raw data, with no meta-data). To use this table, you have to remember what the numbers in each column represent. Also, when you want to select columns, you have to count their position within the table. This can become frustrating and prone to bad errors (getting the columns wrong) especially as the number of columns increase. It is also bad for sending to a colleague, because they will find it hard to remember/use the columns properly.

To solve these problems in Gnuastro’s programs/libraries you aren’t limited to using the column’s number, see Selecting table columns. If the columns have names, units, or comments you can also select your columns based on searches/matches in these fields, for example see Table. Also, in this manner, you can’t guide the program reading the table on how to read the numbers. As an example, the first and third columns above can be read as integer types: the first column might be an ID and the third can be the number of pixels an object occupies in an image. So there is no need to read these to columns as a double type (which takes more memory, and is slower).

In the bare-minimum example above, you also can’t use strings of characters, for example the names of filters, or some other identifier that includes non-numerical characters. In the absence of any information, only numbers can be read robustly. Assuming we read columns with non-numerical characters as string, there would still be the problem that the strings might contain space (or any delimiter) character for some rows. So, each ‘word’ in the string will be interpreted as a column and the program will abort with an error that the rows don’t have the same number of columns.

To correct for these limitations, Gnuastro defines the following convention for storing the table meta-data along with the raw data in one plain text file. The format is primarily designed for ease of reading/writing by eye/fingers, but is also structured enough to be read by a program.

When the first non-white character in a line is #, or there are no non-white characters in it, then the line will not be considered as a row of data in the table (this is a pretty standard convention in many programs, and higher level languages). In the former case, the line is interpreted as a comment. If the comment line starts with ‘# Column N:’, then it is assumed to contain information about column N (a number, counting from 1). Comment lines that don’t start with this pattern are ignored and you can use them to include any further information you want to store with the table in the text file. A column information comment is assumed to have the following format:


Any sequence of characters between ‘:’ and ‘[’ will be interpreted as the column name (so it can contain anything except the ‘[’ character). Anything between the ‘]’ and the end of the line is defined as a comment. Within the brackets, anything before the first ‘,’ is the units (physical units, for example km/s, or erg/s), anything before the second ‘,’ is the short type identifier (see below, and Numeric data types). Finally (still within the brackets), any non-white characters after the second ‘,’ are interpreted as the blank value for that column (see Blank pixels). Note that blank values will be stored in the same type as the column, not as a string83.

When a formatting problem occurs (for example you have specified the wrong type code, see below), or the the column was already given meta-data in a previous comment, or the column number is larger than the actual number of columns in the table (the non-commented or empty lines), then the comment information line will be ignored.

When a comment information line can be used, the leading and trailing white space characters will be stripped from all of the elements. For example in this line:

# Column 5:  column name   [km/s,    f32,-99] Redshift as speed

The NAME field will be ‘column name’ and the TYPE field will be ‘f32’. Note how all the white space characters before and after strings are not used, but those in the middle remained. Also, white space characters aren’t mandatory. Hence, in the example above, the BLANK field will be given the value of ‘-99’.

Except for the column number (N), the rest of the fields are optional. Also, the column information comments don’t have to be in order. In other words, the information for column \(N+m\) (\(m>0\)) can be given in a line before column \(N\). Also, you don’t have to specify information for all columns. Those columns that don’t have this information will be interpreted with the default settings (like the case above: values are double precision floating point, and the column has no name, unit, or comment). So these lines are all acceptable for any table (the first one, with nothing but the column number is redundant):

# Column 5:
# Column 1: ID [,i] The Clump ID.
# Column 3: mag_f160w [AB mag, f] Magnitude from the F160W filter

The data type of the column should be specified with one of the following values:

Note that the FITS binary table standard does not define the unsigned int and unsigned long types, so if you want to convert your tables to FITS binary tables, use other types. Also, note that in the FITS ASCII table, there is only one integer type (long). So if you convert a Gnuastro plain text table to a FITS ASCII table with the Table program, the type information for integers will be lost. Conversely if integer types are important for you, you have to manually set them when reading a FITS ASCII table (for example with the Table program when reading/converting into a file, or with the gnuastro/table.h library functions when reading into memory).

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4.5.3 Selecting table columns

At the lowest level, the only defining aspect of a column in a table is its number, or position. But selecting columns purely by number is not very convenient and, especially when the tables are large it can be very frustrating and prone to errors. Hence, table file formats (for example see Recognized table formats) have ways to store additional information about the columns (meta-data). Some of the most common pieces of information about each column are its name, the units of data in the it, and a comment for longer/informal description of the column’s data.

To facilitate research with Gnuastro, you can select columns by matching, or searching in these three fields, besides the low-level column number. To view the full list of information on the columns in the table, you can use the Table program (see Table) with the command below (replace table-file with the filename of your table, if its FITS, you might also need to specify the HDU/extension which contains the table):

$ asttable --information table-file

Gnuastro’s programs need the columns for different purposes, for example in Crop, you specify the columns containing the central coordinates of the crop centers with the --coordcol option (see Crop options). On the other hand, in MakeProfiles, to specify the column containing the profile position angles, you must use the --pcol option (see MakeProfiles catalog). Thus, there can be no unified common option name to select columns for all programs (different columns have different purposes). However, when the program expects a column for a specific context, the option names end in the col suffix like the examples above. These options accept values in integer (column number), or string (metadata match/search) format.

If the value can be parsed as a positive integer, it will be seen as the low-level column number. Note that column counting starts from 1, so if you ask for column 0, the respective program will abort with an error. When the value can’t be interpreted as an a integer number, it will be seen as a string of characters which will be used to match/search in the table’s meta-data. The meta-data field which the value will be compared with can be selected through the --searchin option, see Input/Output options. --searchin can take three values: name, unit, comment. The matching will be done following this convention:

Note that in both cases, you can ignore the case of alphabetic characters with the --ignorecase option, see Input/Output options. Also, in both cases, multiple columns may be selected with one call to this function. In this case, the order of the selected columns (with one call) will be the same order as they appear in the table.

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4.6 Tessellation

It is sometimes necessary to classify the elements in a dataset (for example pixels in an image) into a grid of individual, non-overlapping tiles. For example when background sky gradients are present in an image, you can define a tile grid over the image. When the tile sizes are set properly, the background’s variation over each tile will be negligible, allowing you to measure (and subtract) it. In other cases (for example spatial domain convolution in Gnuastro, see Convolve), it might simply be for speed of processing: each tile can be processed independently on a separate CPU thread. In the arts and mathematics, this process is formally known as tessellation.

The size of the regular tiles (in units of data-elements, or pixels in an image) can be defined with the --tilesize option. It takes multiple numbers (separated by a comma) which will be the length along the respective dimension (in FORTRAN/FITS dimension order). Divisions are also acceptable, but must result in an integer. For example --tilesize=30,40 can be used for an image (a 2D dataset). The regular tile size along the first FITS axis (horizontal when viewed in SAO ds9) will be 30 pixels and along the second it will be 40 pixels. Ideally, --tilesize should be selected such that all tiles in the image have exactly the same size. In other words, that the dataset length in each dimension is divisible by the tile size in that dimension.

However, this is not always possible: the dataset can be any size and every pixel in it is valuable. In such cases, Gnuastro will look at the significance of the remainder length, if it is not significant (for example one or two pixels), then it will just increase the size of the first tile in the respective dimension and allow the rest of the tiles to have the required size. When the remainder is significant (for example one pixel less than the size along that dimension), the remainder will be added to one regular tile’s size and the large tile will be cut in half and put in the two ends of the grid/tessellation. In this way, all the tiles in the central regions of the dataset will have the regular tile sizes and the tiles on the edge will be slightly larger/smaller depending on the remainder significance. The fraction which defines the remainder significance along all dimensions can be set through --remainderfrac.

The best tile size is directly related to the spatial properties of the property you want to study (for example, gradient on the image). In practice we assume that the gradient is not present over each tile. So if there is a strong gradient (for example in long wavelength ground based images) or the image is of a crowded area where there isn’t too much blank area, you have to choose a smaller tile size. A larger mesh will give more pixels and and so the scatter in the results will be less (better statistics).

For raw image processing, a single tessellation/grid is not sufficient. Raw images are the unprocessed outputs of the camera detectors. Modern detectors usually have multiple readout channels each with its own amplifier. For example the Hubble Space Telescope Advanced Camera for Surveys (ACS) has four amplifiers over its full detector area dividing the square field of view to four smaller squares. Ground based image detectors are not exempt, for example each CCD of Subaru Telescope’s Hyper Suprime-Cam camera (which has 104 CCDs) has four amplifiers, but they have the same height of the CCD and divide the width by four parts.

The bias current on each amplifier is different, and initial bias subtraction is not perfect. So even after subtracting the measured bias current, you can usually still identify the boundaries of different amplifiers by eye. See Figure 11(a) in Akhlaghi and Ichikawa (2015) for an example. This results in the final reduced data to have non-uniform amplifier-shaped regions with higher or lower background flux values. Such systematic biases will then propagate to all subsequent measurements we do on the data (for example photometry and subsequent stellar mass and star formation rate measurements in the case of galaxies).

Therefore an accurate analysis requires a two layer tessellation: the top layer contains larger tiles, each covering one amplifier channel. For clarity we’ll call these larger tiles “channels”. The number of channels along each dimension is defined through the --numchannels. Each channel is then covered by its own individual smaller tessellation (with tile sizes determined by the --tilesize option). This will allow independent analysis of two adjacent pixels from different channels if necessary. If the image is processed or the detector only has one amplifier, you can set the number of channels in both dimension to 1.

The final tessellation can be inspected on the image with the --checktiles option that is available to all programs which use tessellation for localized operations. When this option is called, a FITS file with a _tiled.fits suffix will be created along with the outputs, see Automatic output. Each pixel in this image has the number of the tile that covers it. If the number of channels in any dimension are larger than unity, you will notice that the tile IDs are defined such that the first channels is covered first, then the second and so on. For the full list of processing-related common options (including tessellation options), please see Processing options.

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4.7 Getting help

Probably the first time you read this book, it is either in the PDF or HTML formats. These two formats are very convenient for when you are not actually working, but when you are only reading. Later on, when you start to use the programs and you are deep in the middle of your work, some of the details will inevitably be forgotten. Going to find the PDF file (printed or digital) or the HTML webpage is a major distraction.

GNU software have a very unique set of tools for aiding your memory on the command-line, where you are working, depending how much of it you need to remember. In the past, such command-line help was known as “online” help, because they were literally provided to you ‘on’ the command ‘line’. However, nowadays the word “online” refers to something on the internet, so that term will not be used. With this type of help, you can resume your exciting research without taking your hands off the keyboard.

Another major advantage of such command-line based help routines is that they are installed with the software in your computer, therefore they are always in sync with the executable you are actually running. Three of them are actually part of the executable. You don’t have to worry about the version of the book or program. If you rely on external help (a PDF in your personal print or digital archive or HTML from the official webpage) you have to check to see if their versions fit with your installed program.

If you only need to remember the short or long names of the options, --usage is advised. If it is what the options do, then --help is a great tool. Man pages are also provided for those who are use to this older system of documentation. This full book is also available to you on the command-line in Info format. If none of these seems to resolve the problems, there is a mailing list which enables you to get in touch with experienced Gnuastro users. In the subsections below each of these methods are reviewed.

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4.7.1 --usage

If you give this option, the program will not run. It will only print a very concise message showing the options and arguments. Everything within square brackets ([]) is optional. For example here are the first and last two lines of Crop’s --usage is shown:

$ astcrop --usage
Usage: astcrop [-Do?IPqSVW] [-d INT] [-h INT] [-r INT] [-w INT]
            [-x INT] [-y INT] [-c INT] [-p STR] [-N INT] [--deccol=INT]
            [--setusrconf] [--usage] [--version] [--wcsmode]
            [ASCIIcatalog] FITSimage(s).fits

There are no explanations on the options, just their short and long names shown separately. After the program name, the short format of all the options that don’t require a value (on/off options) is displayed. Those that do require a value then follow in separate brackets, each displaying the format of the input they want, see Options. Since all options are optional, they are shown in square brackets, but arguments can also be optional. For example in this example, a catalog name is optional and is only required in some modes. This is a standard method of displaying optional arguments for all GNU software.

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4.7.2 --help

If the command-line includes this option, the program will not be run. It will print a complete list of all available options along with a short explanation. The options are also grouped by their context. Within each context, the options are sorted alphabetically. Since the options are shown in detail afterwards, the first line of the --help output shows the arguments and if they are optional or not, similar to --usage.

In the --help output of all programs in Gnuastro, the options for each program are classified based on context. The first two contexts are always options to do with the input and output respectively. For example input image extensions or supplementary input files for the inputs. The last class of options is also fixed in all of Gnuastro, it shows operating mode options. Most of these options are already explained in Operating mode options.

The help message will sometimes be longer than the vertical size of your terminal. If you are using a graphical user interface terminal emulator, you can scroll the terminal with your mouse, but we promised no mice distractions! So here are some suggestions:

In case you have a special keyword you are looking for in the help, you don’t have to go through the full list. GNU Grep is made for this job. For example if you only want the list of options whose --help output contains the word “axis” in Crop, you can run the following command:

$ astcrop --help | grep axis

If the output of this option does not fit nicely within the confines of your terminal, GNU does enable you to customize its output through the environment variable ARGP_HELP_FMT, you can set various parameters which specify the formatting of the help messages. For example if your terminals are wider than 70 spaces (say 100) and you feel there is too much empty space between the long options and the short explanation, you can change these formats by giving values to this environment variable before running the program with the --help output. You can define this environment variable in this manner:

$ export ARGP_HELP_FMT=rmargin=100,opt-doc-col=20

This will affect all GNU programs using GNU C library’s argp.h facilities as long as the environment variable is in memory. You can see the full list of these formatting parameters in the “Argp User Customization” part of the GNU C library manual. If you are more comfortable to read the --help outputs of all GNU software in your customized format, you can add your customization (similar to the line above, without the $ sign) to your ~/.bashrc file. This is a standard option for all GNU software.

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4.7.3 Man pages

Man pages were the Unix method of providing command-line documentation to a program. With GNU Info, see Info the usage of this method of documentation is highly discouraged. This is because Info provides a much more easier to navigate and read environment.

However, some operating systems require a man page for packages that are installed and some people are still used to this method of command line help. So the programs in Gnuastro also have Man pages which are automatically generated from the outputs of --version and --help using the GNU help2man program. So if you run

$ man programname

You will be provided with a man page listing the options in the standard manner.

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4.7.4 Info

Info is the standard documentation format for all GNU software. It is a very useful command-line document viewing format, fully equipped with links between the various pages and menus and search capabilities. As explained before, the best thing about it is that it is available for you the moment you need to refresh your memory on any command-line tool in the middle of your work without having to take your hands off the keyboard. This complete book is available in Info format and can be accessed from anywhere on the command-line.

To open the Info format of any installed programs or library on your system which has an Info format book, you can simply run the command below (change executablename to the executable name of the program or library):

$ info executablename

In case you are not already familiar with it, run $ info info. It does a fantastic job in explaining all its capabilities its self. It is very short and you will become sufficiently fluent in about half an hour. Since all GNU software documentation is also provided in Info, your whole GNU/Linux life will significantly improve.

Once you’ve become an efficient navigator in Info, you can go to any part of this book or any other GNU software or library manual, no matter how long it is, in a matter of seconds. It also blends nicely with GNU Emacs (a text editor) and you can search manuals while you are writing your document or programs without taking your hands off the keyboard, this is most useful for libraries like the GNU C library. To be able to access all the Info manuals installed in your GNU/Linux within Emacs, type Ctrl-H + i.

To see this whole book from the beginning in Info, you can run

$ info gnuastro

If you run Info with the particular program executable name, for example astcrop or astnoisechisel:

$ info astprogramname

you will be taken to the section titled “Invoking ProgramName” which explains the inputs and outputs along with the command-line options for that program. Finally, if you run Info with the official program name, for example Crop or NoiseChisel:

$ info ProgramName

you will be taken to the top section which introduces the program. Note that in all cases, Info is not case sensitive.

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4.7.5 help-gnuastro mailing list

Gnuastro maintains the help-gnuastro mailing list for users to ask any questions related to Gnuastro. The experienced Gnuastro users and some of its developers are subscribed to this mailing list and your email will be sent to them immediately. However, when contacting this mailing list please have in mind that they are possibly very busy and might not be able to answer immediately.

To ask a question from this mailing list, send a mail to Anyone can view the mailing list archives at It is best that before sending a mail, you search the archives to see if anyone has asked a question similar to yours. If you want to make a suggestion or report a bug, please don’t send a mail to this mailing list. We have other mailing lists and tools for those purposes, see Report a bug or Suggest new feature.

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4.8 Automatic output

All the programs in Gnuastro are designed such that specifying an output file or directory (based on the program context) is optional. The outputs of most programs are automatically found based on the input and what the program does. For example when you are converting a FITS image named FITSimage.fits to a JPEG image, the JPEG image will be saved in FITSimage.jpg.

Another very important part of the automatic output generation is that all the directory information of the input file name is stripped off of it. This feature can be disabled with the --keepinputdir option, see Common options. It is the default because astronomical data are usually very large and organized specially with special file names. In some cases, the user might not have write permissions in those directories. In fact, even if the data is stored on your own computer, it is advised to only grant write permissions to the super user or root. This way, you won’t accidentally delete or modify your valuable data!

Let’s assume that we are working on a report and want to process the FITS images from two projects (ABC and DEF), which are stored in the sub-directories named ABCproject/ and DEFproject/ of our top data directory (/mnt/data). The following shell commands show how one image from the former is first converted to a JPEG image through ConvertType and then the objects from an image in the latter project are detected using NoiseChisel. The text after the # sign are comments (not typed!).

$ pwd                                               # Current location
$ ls                                         # List directory contents
$ ls /mnt/data/ABCproject                                  # Archive 1
ABC01.fits ABC02.fits ABC03.fits
$ ls /mnt/data/DEFproject                                  # Archive 2
DEF01.fits DEF02.fits DEF03.fits
$ astconvertt /mnt/data/ABCproject/ABC02.fits --output=jpg    # Prog 1
$ ls
ABC01.jpg ABC02.jpg
$ astnoisechisel /mnt/data/DEFproject/DEF01.fits              # Prog 2
$ ls
ABC01.jpg ABC02.jpg DEF01_detected.fits

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4.9 Output headers

The output FITS files created by Gnuastro’s programs have some or all of the following standard keywords to keep the basic date and version information of Gnuastro, its dependencies and the pipeline that is using Gnuastro.


The creation time of the FITS file. This date is written directly by CFITSIO and is in UT format.


Git’s commit description from the running directory of Gnuastro’s programs. If the running directory is not version controlled or libgit2 isn’t installed (see Optional dependencies) then this keyword will not be present. The printed value is equivalent to the output of the following command:

git describe --dirty --always

If the running directory contains non-committed work, then the stored value will have a ‘-dirty’ suffix. This can be very helpful to let you know that the data is not ready to be shared with collaborators or submitted to a journal. You should only share results that are produced after all your work is committed (safely stored in the version controlled history and thus reproducible).

At first sight, version control appears to be mainly a tool for software developers. However progress in a scientific research is almost identical to progress in software development: first you have a rough idea that starts with handful of easy steps. But as the first results appear to be promising, you will have to extend, or generalize, it to make it more robust and work in all the situations your research covers, not just your first test samples. Slowly you will find wrong assumptions or bad implementations that need to be fixed (‘bugs’ in software development parlance). Finally, when you submit the research to your collaborators or a journal, many comments and suggestions will come in, and you have to address them.

Software developers have created version control systems precisely for this kind of activity. Each significant moment in the project’s history is called a “commit”, see Version controlled source. A snapshot of the project in each “commit” is safely stored away, so you can revert back to it at a later time, or check changes/progress. This way, you can be sure that your work is reproducible and track the progress and history. With version control, experimentation in the project’s analysis is greatly facilitated, since you can easily revert back if a brainstorm test procedure fails.

One important feature of version control is that the research result (FITS image, table, report or paper) can be stamped with the unique commit information that produced it. This information will enable you to exactly reproduce that same result later, even if you have made changes/progress. For one example of a research paper’s reproduction pipeline, please see the reproduction pipeline of the paper describing NoiseChisel.


The version of CFITSIO used (see CFITSIO).


The version of WCSLIB used (see WCSLIB). Note that older versions of WCSLIB do not report the version internally. So this is only available if you are using more recent WCSLIB versions.


The version of GNU Scientific Library that was used, see GNU Scientific library.


The version of Gnuastro used (see Version numbering).

Here is one example of the last few lines of an example output.

              / Versions and date
DATE    = '...'                / file creation date
COMMIT  = 'v0-8-g547f6eb'      / Commit description in running dir.
CFITSIO = '3.41    '           / CFITSIO version.
WCSLIB  = '5.16    '           / WCSLIB version.
GSL     = '2.3     '           / GNU Scientific Library version.
GNUASTRO= '0.3'                / GNU Astronomy Utilities version.

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5 Data containers

The most low-level and basic property of a dataset is how it is stored. To process, archive and transmit the data, you need a container to store it first. From the start of the computer age, different formats have been defined to store data, optimized for particular applications. One format/container can never be useful for all applications: the storage defines the application and vice-versa. In astronomy, the Flexible Image Transport System (FITS) standard has become the most common format of data storage and transmission. It has many useful features, for example multiple sub-containers (also known as extensions or header data units, HDUs) within one file, or support for tables as well as images. Each HDU can store an independent dataset and its corresponding meta-data. Therefore, Gnuastro has one program (see Fits) specifically designed to manipulate FITS HDUs and the meta-data (header keywords) in each HDU.

Your astronomical research does not just involve data analysis (where the FITS format is very useful). For example you want to demonstrate your raw and processed FITS images or spectra as figures within slides, reports, or papers. The FITS format is not defined for such applications. Thus, Gnuastro also comes with the ConvertType program (see ConvertType) which can be used to convert a FITS image to and from (where possible) other formats like plain text and JPEG (which allow two way conversion), along with EPS and PDF (which can only be created from FITS, not the other way round).

Finally, the FITS format is not just for images, it can also store tables. Binary tables in particular can be very efficient in storing catalogs that have more than a few tens of columns and rows. However, unlike images (where all elements/pixels have one data type), tables contain multiple columns and each column can have different properties: independent data types (see Numeric data types) and meta-data. In practice, each column can be viewed as a separate container that is grouped with others in the table. The only shared property of the columns in a table is thus the number of elements they contain. To allow easy inspection/manipulation of table columns, Gnuastro has the Table program (see Table). It can be used to select certain table columns in a FITS table and see them as a human readable output on the command-line, or to save them into another plain text or FITS table.

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5.1 Fits

The “Flexible Image Transport System”, or FITS, is by far the most common data container format in astronomy and in constant use since the 1970s. Archiving (future usage, simplicity) has been one of the primary design principles of this format. In the last few decades it has proved so useful and robust that the Vatican Library has also chosen FITS for its “long-term digital preservation” project84.

Although the full name of the standard invokes the idea that it is only for images, it also contains complete and robust features for tables. It started off in the 1970s and was formally published as a standard in 1981, it was adopted by the International Astronomical Union (IAU) in 1982 and an IAU working group to maintain its future was defined in 1988. The FITS 2.0 and 3.0 standards were approved in 2000 and 2008 respectively, and the 4.0 draft has also been released recently, please see the FITS standard document webpage for the full text of all versions. Also see the FITS 3.0 standard paper for a nice introduction and history along with the full standard.

Other formats, for example a JPEG image, only have one image/dataset per file, however one great advantage of the FITS standard is that it allows you to keep multiple datasets (images or tables along with their meta-data) in one file. Each data + metadata is known as an extension, or more formally a header data unit or HDU, in the FITS standard. In theory the HDUs can be completely independent: you can have multiple images of different dimensions/sizes or tables as separate extensions in one file. However, while the standard doesn’t impose any constraints on the relation between the datasets, it is strongly encouraged to group data that are contextually related with each other in one file. For example an image and the table/catalog of objects and their measured properties in that image. Another example can be multiple images of one patch of sky in different colors (filters).

As discussed above, the extensions in a FITS file can be completely independent. To keep some information (meta-data) about the group of extensions in the FITS file, the community has adopted the following convention: put no data in the first extension, so it is just meta-data. This extension can thus be used to store Meta-data regarding the whole file (grouping of extensions). Subsequent extensions may contain data along with their own separate meta-data. All of Gnuastro’s programs also follow this convention: the main dataset (image or table) is in the second extension. See the example list of extension properties in Invoking Fits.

The meta-data contain information about the data, for example which region of the sky an image corresponds to, the units of the data, what telescope, camera, and filter the data were taken with, the software that produced it, or it observation or processing date. Hence without the meta-data, the raw dataset is practically just a collection of numbers and really hard to understand, or connect with the real world (other datasets). It is thus strongly encouraged to supplement your data (at any level of processing) with as much meta-data about your processing/science as possible.

The meta-data of a FITS file is in ASCII format, which can be easily viewed or edited with a text editor or on the command-line. Each meta-data element (known as a keyword generally) is composed of a name, value, units and comments (the last two are optional). For example below you can see three FITS meta-data keywords for specifying the world coordinate system (WCS, or its location in the sky) of a dataset:

LATPOLE =           -27.805089 / [deg] Native latitude of celestial pole
RADESYS = 'FK5'                / Equatorial coordinate system
EQUINOX =               2000.0 / [yr] Equinox of equatorial coordinates

However, there are some limitations which discourage viewing/editing the keywords with text editors. For example there is a fixed length of 80 characters for each keyword (its name, value, units and comments) and there are no new-line characters, so on a text editor all the keywords are seen in one line. Also, the meta-data keywords are immediately followed by the data which are commonly in binary format and will show up as strange looking characters on a text editor, and significantly slowing down the processor.

Gnuastro’s Fits program was designed to allow easy manipulation of FITS extensions and meta-data keywords on the command-line while conforming fully with the FITS standard. For example you can copy or cut (copy and remove) HDUs/extensions from one FITS file to another, or completely delete them. It also has features to delete, add, or edit meta-data keywords within one HDU.

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5.1.1 Invoking Fits

Fits can print or manipulate the FITS file HDUs (extensions), meta-data keywords in a given HDU. The executable name is astfits with the following general template

$ astfits [OPTION...] ASTRdata

One line examples:

## View general information about every extension:
$ astfits image.fits

## Print the header keywords in the second HDU (counting from 0):
$ astfits image.fits -h1

## Only print header keywords that contain `NAXIS':
$ astfits image.fits -h1 | grep NAXIS

## Only print the WCS standard PC matrix elements
$ astfits image.fits -h1 | grep 'PC._.'

## Copy a HDU from input.fits to out.fits:
$ astfits input.fits --copy=hdu-name --output=out.fits

## Update the OLDKEY keyword value to 153.034:
$ astfits --update=OLDKEY,153.034,"Old keyword comment"

## Delete one COMMENT keyword and add a new one:
$ astfits --delete=COMMENT --comment="Anything you like ;-)."

## Write two new keywords with different values and comments:
$ astfits --write=MYKEY1,20.00,"An example keyword" --write=MYKEY2,fd

When no action is requested (and only a file name is given), Fits will print a list of information about the extension(s) in the file. This information includes the HDU number, HDU name (EXTNAME keyword), type of data (see Numeric data types, and the number of data elements it contains (size along each dimension for images and table rows and columns). You can use this to get a general idea of the contents of the FITS file and what HDU to use for further processing, either with the Fits program or any other Gnuastro program.

Here is one example of information about a FITS file with four extensions: the first extension has no data, it is a purely meta-data HDU (commonly used to keep meta-data about the whole file, or grouping of extensions, see Fits). The second extension is an image with name IMAGE and single precision floating point type (float32, see Numeric data types), it has 4287 pixels along its first (horizontal) axis and 4286 pixels along its second (vertical) axis. The third extension is also an image with name MASK. It is in 2-byte integer format (int16) which is commonly used to keep information about pixels (for example to identify which ones were saturated, or which ones had cosmic rays and so on), note how it has the same size as the IMAGE extension. The third extension is a binary table called CATALOG which has 12371 rows and 5 columns (it probably contains information about the sources in the image).

GNU Astronomy Utilities X.X
Run on Day Month DD HH:MM:SS YYYY
HDU (extension) information: `image.fits'.
 Column 1: Index (counting from 0).
 Column 2: Name (`EXTNAME' in FITS standard).
 Column 3: Image data type or `table' format (ASCII or binary).
 Column 4: Size of data in HDU.
0      n/a             uint8           0
1      IMAGE           float32         4287x4286
2      MASK            int16           4287x4286
3      CATALOG         table_binary    12371x5

If a specific HDU is identified on the command-line with the --hdu (or -h option) and no operation requested, then the full list of header keywords in that HDU will be printed (as if the --printallkeys was called, see below). It is important to remember that this only occurs when --hdu is given on the command-line. The --hdu value given in a configuration file will only be used when a specific operation on keywords requested. Therefore as described in the paragraphs above, when no explicit call to the --hdu option is made on the command-line and no operation is requested (on the command-line or configuration files), the basic information of each HDU/extension is printed.

The operating mode and input/output options to Fits are similar to the other programs and fully described in Common options. The options particular to Fits can be divided into two groups: 1) those related to modifying HDUs or extensions (see HDU manipulation), and 2) those related to viewing/modifying meta-data keywords (see Keyword manipulation). These two classes of options cannot be called together in one run: you can either work on the extensions or meta-data keywords in any instance of Fits.

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Each header data unit, or HDU (also known as an extension), in a FITS file is an independent dataset (data + meta-data). Multiple HDUs can be stored in one FITS file, see Fits. The HDU modifying options to the Fits program are listed below.

These options may be called multiple times in one run. If so, the extensions will be copied from the input FITS file to the output FITS file in the given order (on the command-line and also in configuration files, see Configuration file precedence). If the separate classes are called together in one run of Fits, then first --copy is run (on all specified HDUs), followed by --cut (again on all specified HDUs), and then --remove (on all specified HDUs).

The --copy and --cut options need an output FITS file (specified with the --output option). If the output file exists, then the specified HDU will be copied following the last extension of the output file (the existing HDUs in it will be untouched). Thus, after Fits finishes, the copied HDU will be the last HDU of the output file. If no output file name is given, then automatic output will be used to store the HDUs given to this option (see Automatic output).


Copy the specified extension into the output file, see explanations above.

-k STR

Cut (copy to output, remove from input) the specified extension into the output file, see explanations above.


Remove the specified HDU from the input file. From CFITSIO: “In the case of deleting the primary array (the first HDU in the file) then [it] will be replaced by a null primary array containing the minimum set of required keywords and no data.”. So in practice, any existing data (array) and meta-data in the first extension will be removed, but the number of extensions in the file won’t change. This is because of the unique position the first FITS extension has in the FITS standard (for example it cannot be used to store tables).

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The meta-data in each header data unit, or HDU (also known as extension, see Fits) is stored as “keyword”s. Each keyword consists of a name, value, unit, and comments. The Fits program (see Fits) options related to viewing and manipulating keywords in a FITS HDU are described below.

To see the full list of keywords in a FITS HDU, you can use the --printallkeys option. If any of the keywords are to be modified, the headers of the input file will be changed. If you want to keep the original FITS file or HDU, it is easiest to create a copy first and then run Fits on that. In the FITS standard, keywords are always uppercase. So case does not matter in the input or output keyword names you specify.

Most of the options can accept multiple instances in one command. For example you can add multiple keywords to delete by calling --delete multiple times, since repeated keywords are allowed, you can even delete the same keyword multiple times. The action of such options will start from the top most keyword.

The precedence of operations are described below. Note that while the order within each class of actions is preserved, the order of individual actions is not. So irrespective of what order you called --delete and --update. First, all the delete operations are going to take effect then the update operations.

  1. --delete
  2. --rename
  3. --update
  4. --write
  5. --asis
  6. --history
  7. --comment
  8. --date
  9. --printallkeys

All possible syntax errors will be reported before the keywords are actually written. FITS errors during any of these actions will be reported, but Fits won’t stop until all the operations are complete. If --quitonerror is called, then Fits will immediately stop upon the first error.

If you want to inspect only a certain set of header keywords, it is easiest to pipe the output of the Fits program to GNU Grep. Grep is a very powerful and advanced tool to search strings which is precisely made for such situations. For example if you only want to check the size of an image FITS HDU, you can run:

$ astfits input.fits | grep NAXIS

FITS STANDARD KEYWORDS: Some header keywords are necessary for later operations on a FITS file, for example BITPIX or NAXIS, see the FITS standard for their full list. If you modify (for example remove or rename) such keywords, the FITS file extension might not be usable any more. Also be careful for the world coordinate system keywords, if you modify or change their values, any future world coordinate system (like RA and Dec) measurements on the image will also change.

The keyword related options to the Fits program are fully described below.

-a STR

Write STR exactly into the FITS file header with no modifications. If it does not conform to the FITS standards, then it might cause trouble, so please be very careful with this option. If you want to define the keyword from scratch, it is best to use the --write option (see below) and let CFITSIO worry about the standards. The best way to use this option is when you want to add a keyword from one FITS file to another unchanged and untouched. Below is an example of such a case that can be very useful sometimes (on the command-line or in scripts):

$ key=$(astfits firstimage.fits | grep KEYWORD)
$ astfits --asis="$key" secondimage.fits

In particular note the double quotation signs (") around the reference to the key shell variable ($key), since FITS keywords usually have lots of space characters, if this variable is not quoted, the shell will only give the first word in the full keyword to this option, which will definitely be a non-standard FITS keyword and will make it hard to work on the file afterwords. See the “Quoting” section of the GNU Bash manual for more information if your keyword has the special characters $, `, or \.

-d STR

Delete one instance of the STR keyword from the FITS header. Multiple instances of --delete can be given (possibly even for the same keyword, when its repeated in the meta-data). All keywords given will be removed from the headers in the same given order. If the keyword doesn’t exist, Fits will give a warning and return with a non-zero value, but will not stop. To stop as soon as an error occurs, run with --quitonerror.

-r STR

Rename a keyword to a new value. STR contains both the existing and new names, which should be separated by either a comma (,) or a space character. Note that if you use a space character, you have to put the value to this option within double quotation marks (") so the space character is not interpreted as an option separator. Multiple instances of --rename can be given in one command. The keywords will be renamed in the specified order. If the keyword doesn’t exist, Fits will give a warning and return with a non-zero value, but will not stop. To stop as soon as an error occurs, run with --quitonerror.

-u STR

Update a keyword, its value, its comments and its units in the format described below. If there are multiple instances of the keyword in the header, they will be changed from top to bottom (with multiple --update options).

The format of the values to this option can best be specified with an example:

--update=KEYWORD,value,"comments for this keyword",unit

If there is a writing error, Fits will give a warning and return with a non-zero value, but will not stop. To stop as soon as an error occurs, run with --quitonerror.

The value can be any numerical or string value85. Other than the KEYWORD, all the other values are optional. To leave a given token empty, follow the preceding comma (,) immediately with the next. If any space character is present around the commas, it will be considered part of the respective token. So if more than one token has space characters within it, the safest method to specify a value to this option is to put double quotation marks around each individual token that needs it. Note that without double quotation marks, space characters will be seen as option separators and can lead to undefined behavior.

-w STR

Write a keyword to the header. For the possible value input formats, comments and units for the keyword, see the --update option above.

--history STR

Add a HISTORY keyword to the header with the given value. A new HISTORY keyword will be created for every instance of this option. If the string given to this option is longer than 70 characters, it will be separated into multiple keyword cards. If there is an error, Fits will give a warning and return with a non-zero value, but will not stop. To stop as soon as an error occurs, run with --quitonerror.

-c STR
--comment STR

Add a COMMENT keyword to the header with the given value. Similar to the explanation for --history above.


Put the current date and time in the header. If the DATE keyword already exists in the header, it will be updated. If there is a writing error, Fits will give a warning and return with a non-zero value, but will not stop. To stop as soon as an error occurs, run with --quitonerror.


Print all the keywords in the specified FITS extension (HDU) on the command-line. If this option is called along with any of the other keyword editing commands, as described above, all other editing commands take precedence to this. Therefore, it will print the final keywords after all the editing has been done.


Quit if any of the operations above are not successful. By default if an error occurs, Fits will warn the user of the faulty keyword and continue with the rest of actions.

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5.2 ConvertType

The formats of astronomical data were defined mainly for archiving and processing. In other situations, the data might be useful in other formats. For example, when you are writing a paper or report or if you are making slides for a talk, you can’t use a FITS image. Other image formats should be used. In other cases you might want your pixel values in a table format as plain text for input to other programs that don’t recognize FITS, or to include as a table in your report. ConvertType is created for such situations. The various types will increase with future updates and based on need.

The conversion is not only one way (from FITS to other formats), but two ways (except the EPS and PDF formats). So you can convert a JPEG image or text file into a FITS image. Basically, other than EPS, you can use any of the recognized formats as different color channel inputs to get any of the recognized outputs. So before explaining the options and arguments, first a short description of the recognized files types will be given followed a short introduction to digital color.

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5.2.1 Recognized file formats

The various standards and the file name extensions recognized by ConvertType are listed below.


Astronomical data are commonly stored in the FITS format (and in older data sets in IRAF .imh format), a list of file name suffixes which indicate that the file is in this format is given in Arguments.

Each extension of a FITS image only has one value per pixel, so when used as input, each input FITS image contributes as one color channel. If you want multiple extensions in one FITS file for different color channels, you have to repeat the file name multiple times and use the --hdu, --hdu2, --hdu3 or --hdu4 options to specify the different extensions.


The JPEG standard was created by the Joint photographic experts group. It is currently one of the most commonly used image formats. Its major advantage is the compression algorithm that is defined by the standard. Like the FITS standard, this is a raster graphics format, which means that it is pixelated.

A JPEG file can have 1 (for gray-scale), 3 (for RGB) and 4 (for CMYK) color channels. If you only want to convert one JPEG image into other formats, there is no problem, however, if you want to use it in combination with other input files, make sure that the final number of color channels does not exceed four. If it does, then ConvertType will abort and notify you.

The file name endings that are recognized as a JPEG file for input are: .jpg, .JPG, .jpeg, .JPEG, .jpe, .jif, .jfif and .jfi.


TIFF (or Tagged Image File Format) was originally designed as a common format for scanners in the early 90s and since then it has grown to become very general. In many aspects, the TIFF standard is similar to the FITS image standard: it can allow data of many types (see Numeric data types), and also allows multiple images to be stored in a single file (each image in the file is called a ‘directory’ in the TIFF standard). However, unlike FITS, it can only store images, it has no constructs for tables. Another (inconvenient) difference with the FITS standard is that keyword names are stored as numbers, not human-readable text.

However, outside of astronomy, because of its support of different numeric data types, many fields use TIFF images for accurate (for example 16-bit integer or floating point for example) imaging data.

Currently ConvertType can only read TIFF images, if you are interested in writing TIFF images, please get in touch with us.


The Encapsulated PostScript (EPS) format is essentially a one page PostScript file which has a specified size. PostScript also includes non-image data, for example lines and texts. It is a fully functional programming language to describe a document. Therefore in ConvertType, EPS is only an output format and cannot be used as input. Contrary to the FITS or JPEG formats, PostScript is not a raster format, but is categorized as vector graphics.

The Portable Document Format (PDF) is currently the most common format for documents. Some believe that PDF has replaced PostScript and that PostScript is now obsolete. This view is wrong, a PostScript file is an actual plain text file that can be edited like any program source with any text editor. To be able to display its programmed content or print, it needs to pass through a processor or compiler. A PDF file can be thought of as the processed output of the compiler on an input PostScript file. PostScript, EPS and PDF were created and are registered by Adobe Systems.

With these features in mind, you can see that when you are compiling a document with TeX or LaTeX, using an EPS file is much more low level than a JPEG and thus you have much greater control and therefore quality. Since it also includes vector graphic lines we also use such lines to make a thin border around the image to make its appearance in the document much better. No matter the resolution of the display or printer, these lines will always be clear and not pixelated. In the future, addition of text might be included (for example labels or object IDs) on the EPS output. However, this can be done better with tools within TeX or LaTeX such as PGF/Tikz86.

If the final input image (possibly after all operations on the flux explained below) is a binary image or only has two colors of black and white (in segmentation maps for example), then PostScript has another great advantage compared to other formats. It allows for 1 bit pixels (pixels with a value of 0 or 1), this can decrease the output file size by 8 times. So if a gray-scale image is binary, ConvertType will exploit this property in the EPS and PDF (see below) outputs.

The standard formats for an EPS file are .eps, .EPS, .epsf and .epsi. The EPS outputs of ConvertType have the .eps suffix.


As explained above, a PDF document is a static document description format, viewing its result is therefore much faster and more efficient than PostScript. To create a PDF output, ConvertType will make a PostScript page description and convert that to PDF using GPL Ghostscript. The suffixes recognized for a PDF file are: .pdf, .PDF. If GPL Ghostscript cannot be run on the PostScript file, it will remain and a warning will be printed.


This is not actually a file type! But can be used to fill one color channel with a blank value. If this argument is given for any color channel, that channel will not be used in the output.

Plain text

Plain text files have the advantage that they can be viewed with any text editor or on the command-line. Most programs also support input as plain text files. As input, each plain text file is considered to contain one color channel.

In ConvertType, the recognized extensions for plain text files are .txt and .dat. As described in Invoking ConvertType, if you just give these extensions, (and not a full filename) as output, then automatic output will be preformed to determine the final output name (see Automatic output). Besides these, when the format of a file cannot be recognized from its name, ConvertType will fall back to plain text mode. So you can use any name (even without an extension) for a plain text input or output. Just note that when the suffix is not recognized, automatic output will not be preformed.

The basic input/output on plain text images is very similar to how tables are read/written as described in Gnuastro text table format. Simply put, the restrictions are very loose, and there is a convention to define a name, units, data type (see Numeric data types), and comments for the data in a commented line. The only difference is that as a table, a text file can contain many datasets (columns), but as a 2D image, it can only contain one dataset. As a result, only one information comment line is necessary for a 2D image, and instead of the starting ‘# Column N’ (N is the column number), the information line for a 2D image must start with ‘# Image 1’. When ConvertType is asked to output to plain text file, this information comment line is written before the image pixel values.

When converting an image to plain text, consider the fact that if the image is large, the number of columns in each line will become very large, possibly making it very hard to open in some text editors.

Standard output (command-line)

This is very similar to the plain text output, but instead of creating a file to keep the printed values, they are printed on the command line. This can be very useful when you want to redirect the results directly to another program in one command with no intermediate file. The only difference is that only the pixel values are printed (with no information comment line). To print to the standard output, set the output name to ‘stdout’.

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5.2.2 Color

An image is a two dimensional array of 2 dimensional elements called pixels. If each pixel only has one value, the image is known as a gray-scale image and no color is defined. The range of values in the image can be interpreted as shades of any color, it is customary to use shades of black or gray-scale. However, to produce the color spectrum in the digital world, several primary colors must be mixed. Therefore in a color image, each pixel has several values depending on how many primary colors were chosen. For example on the digital monitor or color digital cameras, all colors are built by mixing the three colors of Red-Green-Blue (RGB) with various proportions. However, for printing on paper, standard printers use the Cyan-Magenta-Yellow-Key (CMYK, Key=black) color space. Therefore when printing an RGB image, usually a transformation of color spaces will be necessary.

In a colored digital camera, a color image is produced by dividing the pixel’s area between three colors (filters). However in astronomy due to the intrinsic faintness of most of the targets, the collecting area of the pixel is very important for us. Hence the full area of the pixel is used and one value is stored for that pixel in the end. One color filter is used for the whole image. Thus a FITS image is inherently a gray-scale image and no color can be defined for it.

One way to represent a gray-scale image in different color spaces is to use the same proportions of the primary colors in each pixel. This is the common way most FITS image converters work: they fill all the channels with the same values. The downside is two fold:

To solve both these problems, the best way is to save the FITS image into the black channel of the CMYK color space. In the RGB color space all three channels have to be used. The JPEG standard is the only common standard that accepts CMYK color space, that is why currently only the JPEG standard is included and not the PNG standard for example.

The JPEG and EPS standards set two sizes for the number of bits in each channel: 8-bit and 12-bit. The former is by far the most common and is what is used in ConvertType. Therefore, each channel should have values between 0 to 2^8-1=255. From this we see how each pixel in a gray-scale image is one byte (8 bits) long, in an RGB image, it is 3bytes long and in CMYK it is 4bytes long. But thanks to the JPEG compression algorithms, when all the pixels of one channel have the same value, that channel is compressed to one pixel. Therefore a Grayscale image and a CMYK image that has only the K-channel filled are approximately the same file size.

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5.2.3 Invoking ConvertType

ConvertType will convert any recognized input file type to any specified output type. The executable name is astconvertt with the following general template

$ astconvertt [OPTION...] InputFile [InputFile2] ... [InputFile4]

One line examples:

## Convert an image in FITS to PDF:
$ astconvertt image.fits --output=pdf

## Convert an image in JPEG to FITS:
$ astconvertt image.jpg -ogalaxy.fits

## Use three plain text 2D arrays to create an RGB JPEG output:
$ astconvertt f1.txt f2.txt f3.fits -o.jpg

## Use two images and one blank for an RGB EPS output:
$ astconvertt M31_r.fits M31_g.fits blank -oeps

The file type of the output will be specified with the (possibly complete) file name given to the --output option, which can either be given on the command-line or in any of the configuration files (see Configuration files). Note that if the output suffix is not recognized, it will default to plain text format, see Recognized file formats.

The order of multiple input files is important. After reading the input file(s) the number of color channels in all the inputs will be used to define which color space is being used for the outputs and how each color channel is interpreted. Note that one file might have more than one color channel (for example in the JPEG format). If there is one color channel the output is gray-scale, if three input color channels are given they are respectively considered to be the red, green and blue color channels and if there are four color channels they are respectively considered to be cyan, magenta, yellow and black.

The value to --output (or -o) can be either a full file name or just the suffix of the desired output format. In the former case, that same name will be used for the output. In the latter case, the name of the output file will be set based on the automatic output guidelines, see Automatic output. Note that the suffix name can optionally start a . (dot), so for example --output=.jpg and --output=jpg are equivalent. See Recognized file formats

Besides the common set of options explained in Common options, the options to ConvertType can be classified into input, output and flux related options. The majority of the options are to do with the flux range. Astronomical data usually have a very large dynamic range (difference between maximum and minimum value) and different subjects might be better demonstrated with a limited flux range.



In ConvertType, it is possible to call the HDU option multiple times for the different input FITS or TIFF files in the same order that they are called on the command-line. Note that in the TIFF standard, one ‘directory’ (similar to a FITS HDU) may contain multiple color channels (for example when the image is in RGB).

Except for the fact that multiple calls are possible, this option is identical to the common --hdu in Input/Output options. The number of calls to this option cannot be less than the number of input FITS or TIFF files, but if there are more, the extra HDUs will be ignored, note that they will be read in the order described in Configuration file precedence.

Unlike CFITSIO, libtiff (which is used to read TIFF files) only recognizes numbers (counting from zero, similar to CFITSIO) for ‘directory’ identification. Hence the concept of names is not defined for the directories and the values to this option for TIFF files must be numbers.


-w FLT

The width of the output in centimeters. This is only relevant for those formats that accept such a width (not plain text for example). For most digital purposes, the number of pixels is far more important than the value to this parameter because you can adjust the absolute width (in inches or centimeters) in your document preparation program.

-b INT

The width of the border to be put around the EPS and PDF outputs in units of PostScript points. There are 72 or 28.35 PostScript points in an inch or centimeter respectively. In other words, there are roughly 3 PostScript points in every millimeter. If you are planning on adding a border, its significance is highly correlated with the value you give to the --widthincm parameter.

Unfortunately in the document structuring convention of the PostScript language, the “bounding box” has to be in units of PostScript points with no fractions allowed. So the border values only have to be specified in integers. To have a final border that is thinner than one PostScript point in your document, you can ask for a larger width in ConvertType and then scale down the output EPS or PDF file in your document preparation program. For example by setting width in your includegraphics command in TeX or LaTeX. Since it is vector graphics, the changes of size have no effect on the quality of your output quality (pixels don’t get different values).


Use Hexadecimal encoding in creating EPS output. By default the ASCII85 encoding is used which provides a much better compression ratio. When converted to PDF (or included in TeX or LaTeX which is finally saved as a PDF file), an efficient binary encoding is used which is far more efficient than both of them. The choice of EPS encoding will thus have no effect on the final PDF.

So if you want to transfer your EPS files (for example if you want to submit your paper to arXiv or journals in PostScript), their storage might become important if you have large images or lots of small ones. By default ASCII85 encoding is used which offers a much better compression ratio (nearly 40 percent) compared to Hexadecimal encoding.

-u INT

The quality (compression) of the output JPEG file with values from 0 to 100 (inclusive). For other formats the value to this option is ignored. Note that only in gray-scale (when one input color channel is given) will this actually be the exact quality (each pixel will correspond to one input value). If it is in color mode, some degradation will occur. While the JPEG standard does support loss-less graphics, it is not commonly supported.

Flux range:

-c STR

(=STR) Change pixel values with the following format "from1:to1, from2:to2,...". This option is very useful in displaying labeled pixels (not actual data images which have noise) like segmentation maps. In labeled images, usually a group of pixels have a fixed integer value. With this option, you can manipulate the labels before the image is displayed to get a better output for print or to emphasize on a particular set of labels and ignore the rest. The labels in the images will be changed in the same order given. By default first the pixel values will be converted then the pixel values will be truncated (see --fluxlow and --fluxhigh).

You can use any number for the values irrespective of your final output, your given values are stored and used in the double precision floating point format. So for example if your input image has labels from 1 to 20000 and you only want to display those with labels 957 and 11342 then you can run ConvertType with these options:

$ astconvertt --change=957:50000,11342:50001 --fluxlow=5e4 \
   --fluxhigh=1e5 segmentationmap.fits --output=jpg

While the output JPEG format is only 8 bit, this operation is done in an intermediate step which is stored in double precision floating point. The pixel values are converted to 8-bit after all operations on the input fluxes have been complete. By placing the value in double quotes you can use as many spaces as you like for better readability.


Change pixel values (with --change) after truncation of the flux values, by default it is the opposite.


The minimum flux (pixel value) to display in the output image, any pixel value below this value will be set to this value in the output. If the value to this option is the same as --fluxhigh, then no flux truncation will be applied. Note that when multiple channels are given, this value is used for all the color channels.


The maximum flux (pixel value) to display in the output image, see --fluxlow.

-m INT

This is only used for the JPEG and EPS output formats which have an 8-bit space for each channel of each pixel. The maximum value in each pixel can therefore be \(2^8-1=255\). With this option you can change (decrease) the maximum value. By doing so you will decrease the dynamic range. It can be useful if you plan to use those values for other purposes.


If the lowest pixel value in the input channels is larger than the value to --fluxlow, then that input value will be redundant. In some situations it might be necessary to set the minimum byte value (0) to correspond to that flux even if the data do not reach that value. With this option you can do this. Note that if the minimum pixel value is smaller than --fluxlow, then this option is redundant.


See --flminbyte.


For 8-bit output types (JPEG, EPS, and PDF for example) the final value that is stored is inverted so white becomes black and vice versa. The reason for this is that astronomical images usually have a very large area of blank sky in them. The result will be that a large are of the image will be black. Note that this behavior is ideal for gray-scale images, if you want a color image, the colors are going to be mixed up.

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5.3 Table

Tables are the products of processing astronomical images and spectra. For example in Gnuastro, MakeCatalog will process the defined pixels over an object and produce a catalog (see MakeCatalog). For each identified object, MakeCatalog can print its position on the image or sky, its total brightness and many other information that is deducible from the given image. Each one of these properties is a column in its output catalog (or table) and for each input object, we have a row.

When there are only a small number of objects (rows) and not too many properties (columns), then a simple plain text file is mainly enough to store, transfer, or even use the produced data. However, to be more efficient in all these aspects, astronomers have defined the FITS binary table standard to store data in a binary (0 and 1) format, not plain text. This can offer major advantages in all those aspects: the file size will be greatly reduced and the reading and writing will be faster (because the RAM and CPU also work in binary).

The FITS standard also defines a standard for ASCII tables, where the data are stored in the human readable ASCII format, but within the FITS file structure. These are mainly useful for keeping ASCII data along with images and possibly binary data as multiple (conceptually related) extensions within a FITS file. The acceptable table formats are fully described in Tables.

Binary tables are not easily readable by human eyes. There is no fixed/unified standard on how the zero and ones should be interpreted. The Unix-like operating systems have flourished because of a simple fact: communication between the various tools is based on human readable characters87. So while the FITS table standards are very beneficial for the tools that recognize them, they are hard to use in the vast majority of available software. This creates limitations for their generic use.

‘Table’ is Gnuastro’s solution to this problem. With Table, FITS tables (ASCII or binary) are directly accessible to the Unix-like operating systems power-users (those working the command-line or shell, see Command-line interface). With Table, a FITS table (in binary or ASCII formats) is only one command away from AWK (or any other tool you want to use). Just like a plain text file that you read with the cat command. You can pipe the output of Table into any other tool for higher-level processing, see the examples in Invoking Table for some simple examples.

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5.3.1 Invoking Table

Table will read/write, select, convert, or show the information of the columns in FITS ASCII table, FITS binary table and plain text table files, see Tables. Output columns can also be determined by number or regular expression matching of column names, units, or comments. The executable name is asttable with the following general template

$ asttable [OPTION...] InputFile

One line examples:

## Get the table column information (name, data type, or units):
$ asttable bintab.fits --information

## Print columns named RA and DEC, followed by all the columns where
## the name starts with "MAG_":
$ asttable bintab.fits --column=RA --column=DEC --column=/^MAG_/

## Similar to above, but with one call to --column (or -c) and writes
## the columns to a file (with metadata) instead of the command-line.
$ asttable bintab.fits -cRA,DEC,/^MAG_/ --output=out.txt

## Only print the 2nd column, and the third column multiplied by 5:
$ asttable bintab.fits -c2,5 | awk '{print $1, 5*$2}'

## Only print rows with a value in the 10th column above 100000:
$ asttable bintab.fits | awk '$10>10e5 {print}'

## Sort the output columns by the third column, save output:
$ asttable bintab.fits | 'sort -nk3 > output.txt

## Convert a plain text table to a binary FITS table:
$ asttable plaintext.txt --output=table.fits --tabletype=fits-binary

In the absence of selected columns, all the input file’s columns will be output. If the specified output is a FITS file, the type of FITS table (binary or ASCII) will be determined from the --tabletype option. If the output is not a FITS file, it will be printed as a plain text table (with space characters between the columns). When the columns are accompanied by meta-data (like column name, units, or comments), this information will also printed in the plain text file before the table, as described in Gnuastro text table format.

For the full list of options common to all Gnuastro programs please see Common options. Options can also be stored in directory, user or system-wide configuration files to avoid repeating on the command-line, see Configuration files. Table does not follow Automatic output that is common in most Gnuastro programs, see Automatic output. Thus, in the absence of an output file, the selected columns will be printed on the command-line with no column information, ready for redirecting to other tools like AWK or sort, similar to the examples above.


Only print the column information in the specified table on the command-line and exit. Each column’s information (number, name, units, data type, and comments) will be printed as a row on the command-line. Note that the FITS standard only requires the data type (see Numeric data types), and in plain text tables, no meta-data/information is mandatory. Gnuastro has its own convention in the comments of a plain text table to store and transfer this information as described in Gnuastro text table format.

This option will take precedence over the --column option, so when it is called along with requested columns, the latter will be ignored. This can be useful if you forget the identifier of a column after you have already typed some on the command-line. You can simply add a -i and run Table to see the whole list and remember. Then you can use the shell history (with the up arrow key on the keyboard), and retrieve the last command with all the previously typed columns present, delete -i and add the identifier you had forgot.


Specify the columns to read, see Selecting table columns for a thorough explanation on how the value to this option is interpreted. There are two ways to specify multiple columns: 1) multiple calls to this option, 2) using a comma between each column specifier in one call to this option. These different solutions may be mixed in one call to Table: for example, -cRA,DEC -cMAG, or -cRA -cDEC -cMAG are both equivalent to -cRA -cDEC -cMAG. The order of the output columns will be the same order given to the option or in the configuration files (see Configuration file precedence).

This option is not mandatory, if no specific columns are requested, all the input table columns are output. When this option is called multiple times, it is possible to output one column more than once.

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6 Data manipulation

Images are one of the major formats of data that is used in astronomy. The functions in this chapter explain the GNU Astronomy Utilities which are provided for their manipulation. For example cropping out a part of a larger image or convolving the image with a given kernel or applying a transformation to it.

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6.1 Crop

Astronomical images are often very large, filled with thousands of galaxies. It often happens that you only want a section of the image, or you have a catalog of sources and you want to visually analyze them in small postage stamps. Crop is made to do all these things. When more than one crop is required, Crop will divide the crops between multiple threads to significantly reduce the run time.

Astronomical surveys are usually extremely large. So large in fact, that the whole survey will not fit into a reasonably sized file. Because of this, surveys usually cut the final image into separate tiles and store each tile in a file. For example the COSMOS survey’s Hubble space telescope, ACS F814W image consists of 81 separate FITS images, with each one having a volume of 1.7 Giga bytes.

Even though the tile sizes are chosen to be large enough that too many galaxies/targets don’t fall on the edges of the tiles, inevitably some do. So when you simply crop the image of such targets from one tile, you will miss a large area of the surrounding sky (which is essential in estimating the noise). Therefore in its WCS mode, Crop will stitch parts of the tiles that are relevant for a target (with the given width) from all the input images that cover that region into the output. Of course, the tiles have to be present in the list of input files.

Besides cropping postage stamps around certain coordinates, Crop can also crop arbitrary polygons from an image (or a set of tiles by stitching the relevant parts of different tiles within the polygon), see --polygon in Invoking Crop. Alternatively, it can crop out rectangular regions through the --section option from one image, see Crop section syntax.

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6.1.1 Crop modes

In order to be comprehensive, intuitive, and easy to use, there are two ways to define the crop:

  1. From its center and side length. For example if you already know the coordinates of an object and want to inspect it in an image or to generate postage stamps of a catalog containing many such coordinates.
  2. The vertices of the crop region, this can be useful for larger crops over many targets, for example to crop out a uniformly deep, or contiguous, region of a large survey.

Irrespective of how the crop region is defined, the coordinates to define the crop can be in Image (pixel) or World Coordinate System (WCS) standards. All coordinates are read as floating point numbers (not integers, except for the --section option, see below). By setting the mode in Crop, you define the standard that the given coordinates must be interpreted. Here, the different ways to specify the crop region are discussed within each standard. For the full list options, please see Invoking Crop.

When the crop is defined by its center, the respective (integer) central pixel position will be found internally according to the FITS standard. To have this pixel positioned in the center of the cropped region, the final cropped region will have an add number of pixels (even if you give an even number to --width in image mode).

Furthermore, when the crop is defined as by its center, Crop allows you to only keep crops what don’t have any blank pixels in the vicinity of their center (your primary target). This can be very convenient when your input catalog/coordinates originated from another survey/filter which is not fully covered by your input image, to learn more about this feature, please see the description of the --checkcenter option in Invoking Crop.

Image coordinates

In image mode (--mode=img), Crop interprets the pixel coordinates and widths in units of the input data-elements (for example pixels in an image, not world coordinates). In image mode, only one image may be input. The output crop(s) can be defined in multiple ways as listed below.

Center of multiple crops (in a catalog)

The center of (possibly multiple) crops are read from a text file. In this mode, the columns identified with the --coordcol option are interpreted as the center of a crop with a width of --width pixels along each dimension. The columns can contain any floating point value. The value to --output option is seen as a directory which will host (the possibly multiple) separate crop files, see Crop output for more. For a tutorial using this feature, please see Hubble visually checks and classifies his catalog.

Center of a single crop (on the command-line)

The center of the crop is given on the command-line with the --center option. The crop width is specified by the --width option along each dimension. The given coordinates and width can be any floating point number.

Vertices of a single crop

In Image mode there are two options to define the vertices of a region to crop: --section and --polygon. The former is lower-level (doesn’t accept floating point vertices, and only a rectangular region can be defined), it is also only available in Image mode. Please see Crop section syntax for a full description of this method.

The latter option (--polygon) is a higher-level method to define any convex polygon (with any number of vertices) with floating point values. Please see the description of this option in Invoking Crop for its syntax.

WCS coordinates

In WCS mode (--mode=wcs), the coordinates and widths are interpreted using the World Coordinate System (WCS, that must accompany the dataset), not pixel coordinates. In WCS mode, Crop accepts multiple datasets as input. When the cropped region (defined by its center or vertices) overlaps with multiple of the input images/tiles, the overlapping regions will be taken from the respective input (they will be stitched when necessary for each output crop).

In this mode, the input images do not necessarily have to be the same size, they just need to have the same orientation and pixel resolution. Currently only orientation along the celestial coordinates is accepted, if your input has a different orientation you can use Warp’s --align option to align the image before cropping it (see Warp).

Each individual input image/tile can even be smaller than the final crop. In any case, any part of any of the input images which overlaps with the desired region will be used in the crop. Note that if there is an overlap in the input images/tiles, the pixels from the last input image read are going to be used for the overlap. Crop will not change pixel values, so it assumes your overlapping tiles were cutout from the same original image. There are multiple ways to define your cropped region as listed below.

Center of multiple crops (in a catalog)

Similar to catalog inputs in Image mode (above), except that the values along each dimension are assumed to have the same units as the dataset’s WCS information. For example, the central RA and Dec value for each crop will be read from the first and second calls to the --coordcol option. The width of the cropped box (in units of the WCS, or degrees in RA and Dec mode) must be specified with the --width option.

Center of a single crop (on the command-line)

You can specify the center of only one crop box with the --center option. If it exists in the input images, it will be cropped similar to the catalog mode, see above also for --width.

Vertices of a single crop

The --polygon option is a high-level method to define any convex polygon (with any number of vertices). Please see the description of this option in Invoking Crop for its syntax.

CAUTION: In WCS mode, the image has to be aligned with the celestial coordinates, such that the first FITS axis is parallel (opposite direction) to the Right Ascension (RA) and the second FITS axis is parallel to the declination. If these conditions aren’t met for an image, Crop will warn you and abort. You can use Warp’s --align option to align the input image with these coordinates, see Warp.

As a summary, if you don’t specify a catalog, you have to define the cropped region manually on the command-line. In any case the mode is mandatory for Crop to be able to interpret the values given as coordinates or widths.

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6.1.2 Crop section syntax

When in image mode, one of the methods to crop only one rectangular section from the input image is to use the --section option. Crop has a powerful syntax to read the box parameters from a string of characters. If you leave certain parts of the string to be empty, Crop can fill them for you based on the input image sizes.

To define a box, you need the coordinates of two points: the first (X1, Y1) and the last pixel (X2, Y2) pixel positions in the image, or four integer numbers in total. The four coordinates can be specified with one string in this format: ‘X1:X2,Y1:Y2’. This string is given to the --section option. Therefore, the pixels along the first axis that are \(\geq\)X1 and \(\leq\)X2 will be included in the cropped image. The same goes for the second axis. Note that each different term will be read as an integer, not a float. This is a low-level option, for a higher-level way to specify region (any polygon, not just a box), please see the --polygon option in Crop options. Also note that in the FITS standard, pixel indexes along each axis start from unity(1) not zero(0).

You can omit any of the values and they will be filled automatically. The left hand side of the colon (:) will be filled with 1, and the right side with the image size. So, 2:,: will include the full range of pixels along the second axis and only those with a first axis index larger than 2 in the first axis. If the colon is omitted for a dimension, then the full range is automatically used. So the same string is also equal to 2:, or 2: or even 2. If you want such a case for the second axis, you should set it to: ,2.

If you specify a negative value, it will be seen as before the indexes of the image which are outside the image along the bottom or left sides when viewed in SAO ds9. In case you want to count from the top or right sides of the image, you can use an asterisk (*). When confronted with a *, Crop will replace it with the maximum length of the image in that dimension. So *-10:*+10,*-20:*+20 will mean that the crop box will be 20\times40 pixels in size and only include the top corner of the input image with 3/4 of the image being covered by blank pixels, see Blank pixels.

If you feel more comfortable with space characters between the values, you can use as many space characters as you wish, just be careful to put your value in double quotes, for example --section="5:200, 123:854". If you forget the quotes, anything after the first space will not be seen by --section and you will most probably get an error because the rest of your string will be read as a filename (which most probably doesn’t exist). See Command-line for a description of how the command-line works.

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6.1.3 Blank pixels

The cropped box can potentially include pixels that are beyond the image range. For example when a target in the input catalog was very near the edge of the input image. The parts of the cropped image that were not in the input image will be filled with the following two values depending on the data type of the image. In both cases, SAO ds9 will not color code those pixels.

You can ask for such blank regions to not be included in the output crop image using the --noblank option. In such cases, there is no guarantee that the image size of your outputs are what you asked for.

In some survey images, unfortunately they do not use the BLANK FITS keyword. Instead they just give all pixels outside of the survey area a value of zero. So by default, when dealing with float or double image types, any values that are 0.0 are also regarded as blank regions. This can be turned off with the --zeroisnotblank option.

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6.1.4 Invoking Crop

Crop will crop a region from an image. If in WCS mode, it will also stitch parts from separate images in the input files. The executable name is astcrop with the following general template

$ astcrop [OPTION...] [ASCIIcatalog] ASTRdata ...

One line examples:

## Crop all objects in cat.txt from image.fits:
$ astcrop --catalog=cat.txt image.fits

## Crop all options in catalog (with RA,DEC) from all the files
## ending in `_drz.fits' in `/mnt/data/COSMOS/':
$ astcrop --mode=wcs --catalog=cat.txt /mnt/data/COSMOS/*_drz.fits

## Crop the outer 10 border pixels of the input image:
$ astcrop --section=10:*-10,10:*-10 --hdu=2 image.fits

## Crop region around RA and Dec of (189.16704, 62.218203):
$ astcrop --mode=wcs --center=189.16704,62.218203 goodsnorth.fits

## Crop region around pixel coordinate (568.342, 2091.719):
$ astcrop --mode=img --center=568.342,2091.719 --width=201 image.fits

Crop has one mandatory argument which is the input image name(s), shown above with ASTRdata .... You can use shell expansions, for example * for this if you have lots of images in WCS mode. If the crop box centers are in a catalog, you can use the --catalog option. In other cases, you have to provide the single cropped output parameters must be given with command-line options. See Crop output for how the output file name(s) can be specified. For the full list of general options to all Gnuastro programs (including Crop), please see Common options.

Floating point numbers can be used to specify the crop region (except the --section option, see Crop section syntax). In such cases, the floating point values will be used to find the desired integer pixel indices based on the FITS standard. Hence, Crop ultimately doesn’t do any sub-pixel cropping (in other words, it doesn’t change pixel values). If you need such crops, you can use Warp to first warp the image to the a new pixel grid, then crop from that. For example, let’s assume you want a crop from pixels 12.982 to 80.982 along the first dimension. You should first translate the image by \(-0.482\) (note that the edge of a pixel is at integer multiples of \(0.5\)). So you should run Warp with --translate=-0.482,0 and then crop the warped image with --section=13:81.

There are two ways to define the cropped region: with its center or its vertices. See Crop modes for a full description. In the former case, Crop can check if the central region of the cropped image is indeed filled with data or is blank (see Blank pixels), and not produce any output when the center is blank, see the description under --checkcenter for more.

When in catalog mode, Crop will run in parallel unless you set --numthreads=1, see Multi-threaded operations. Note that when multiple outputs are created with threads, the outputs will not be created in the same order. This is because the threads are asynchronous and thus not started in order. This has no effect on each output, see Hubble visually checks and classifies his catalog for a tutorial on effectively using this feature.

Next: , Previous: , Up: Invoking astcrop   [Contents][Index] Crop options

The options can be classified into the following contexts: Input, Output and operating mode options. Options that are common to all Gnuastro program are listed in Common options and will not be repeated here.

When you are specifying the crop vertices your self (through --section, or --polygon) on relatively small regions (depending on the resolution of your images) the outputs from image and WCS mode can be approximately equivalent. However, as the crop sizes get large, the curved nature of the WCS coordinates have to be considered. For example, when using --section, the right ascension of the bottom left and top left corners will not be equal. If you only want regions within a given right ascension, use --polygon in WCS mode.

Input image parameters:


Specify the first keyword card (line number) to start finding the input image world coordinate system information. Distortions were only recently included in WCSLIB (from version 5). Therefore until now, different telescope would apply their own specific set of WCS keywords and put them into the image header along with those that WCSLIB does recognize. So now that WCSLIB recognizes most of the standard distortion parameters, they will get confused with the old ones and give completely wrong results. For example in the CANDELS-GOODS South images88.

The two --hstartwcs and --hendwcs are thus provided so when using older datasets, you can specify what region in the FITS headers you want to use to read the WCS keywords. Note that this is only relevant for reading the WCS information, basic data information like the image size are read separately. These two options will only be considered when the value to --hendwcs is larger than that of --hstartwcs. So if they are equal or --hstartwcs is larger than --hendwcs, then all the input keywords will be parsed to get the WCS information of the image.


Specify the last keyword card to read for specifying the image world coordinate system on the input images. See --hstartwcs

Crop box parameters:

-c FLT[,FLT[,...]]

The central position of the crop in the input image. The positions along each dimension must be separated by a comma (,) and fractions are also acceptable. The number of values given to this option must be the same as the dimensions of the input dataset. The width of the crop should be set with --width. The units of the coordinates are read based on the value to the --mode option, see below.

-w FLT[,FLT[,...]]

Width of the cropped region about its center. --width may take either a single value (to be used for all dimensions) or multiple values (a specific value for each dimension). If in WCS mode, value(s) given to this option will be read in the same units as the dataset’s WCS information along this dimension. The final output will have an odd number of pixels to allow easy identification of the pixel which keeps your requested coordinate (from --center or --catalog).

The --width option also accepts fractions. For example if you want the width of your crop to be 3 by 5 arcseconds along RA and Dec respectively, you can call it with: --width=3/3600,5/3600.

If you want an even sided crop, you can run Crop afterwards with --section=":*-1,:*-1" or --section=2:,2: (depending on which side you don’t need), see Crop section syntax.

-l STR

String of crop polygon vertices. Note that currently only convex polygons should be used. In the future we will make it work for all kinds of polygons. Convex polygons are polygons that do not have an internal angle more than 180 degrees. This option can be used both in the image and WCS modes, see Crop modes. The cropped image will be the size of the rectangular region that completely encompasses the polygon. By default all the pixels that are outside of the polygon will be set as blank values (see Blank pixels). However, if --outpolygon is called all pixels internal to the vertices will be set to blank.

The syntax for the polygon vertices is similar to, and simpler than, that for --section. In short, the dimensions of each coordinate are separated by a comma (,) and each vertex is separated by a colon (:). You can define as many vertices as you like. If you would like to use space characters between the dimensions and vertices to make them more human-readable, then you have to put the value to this option in double quotation marks.

For example, let’s assume you want to work on the deepest part of the WFC3/IR images of Hubble Space Telescope eXtreme Deep Field (HST-XDF). According to the webpage89 the deepest part is contained within the coordinates:

[ (53.187414,-27.779152), (53.159507,-27.759633),
  (53.134517,-27.787144), (53.161906,-27.807208) ]

They have provided mask images with only these pixels in the WFC3/IR images, but what if you also need to work on the same region in the full resolution ACS images? Also what if you want to use the CANDELS data for the shallow region? Running Crop with --polygon will easily pull out this region of the image for you irrespective of the resolution. If you have set the operating mode to WCS mode in your nearest configuration file (see Configuration files), there is no need to call --mode=wcs on the command line. You may also provide many FITS images/tiles and Crop will stitch them to produce this cropped region:

$ astcrop --mode=wcs desired-filter-image(s).fits           \
   --polygon="53.187414,-27.779152 : 53.159507,-27.759633 : \
              53.134517,-27.787144 : 53.161906,-27.807208"

In other cases, you have an image and want to define the polygon yourself (it isn’t already published like the example above). As the number of vertices increases, checking the vertex coordinates on a FITS viewer (for example SAO ds9) and typing them in one by one can be very tedious and prone to typo errors.

You can take the following steps to avoid the frustration and possible typos: Open the image with ds9 and activate its “region” mode with Edit→Region. Then define the region as a polygon with Region→Shape→Polygon. Click on the approximate center of the region you want and a small square will appear. By clicking on the vertices of the square you can shrink or expand it, clicking and dragging anywhere on the edges will enable you to define a new vertex. After the region has been nicely defined, save it as a file with Region→Save Regions. You can then select the name and address of the output file, keep the format as REG and press “OK”. In the next window, keep format as “ds9” and “Coordinate System” as “fk5”. A plain text file (let’s call it ds9.reg) is now created.

You can now convert this plain text file to Crop’s polygon format with this command (when typing on the command-line, ignore the “\” at the end of the first and second lines along with the extra spaces, these are only for nice printing):

$ v=$(awk 'NR==4' ds9.reg | sed -e's/polygon(//'        \
           -e's/\([^,]*,[^,]*\),/\1:/g' -e's/)//' )
$ astcrop --mode=wcs image.fits --polygon=$v

Keep all the regions outside the polygon and mask the inner ones with blank pixels (see Blank pixels). This is practically the inverse of the default mode of treating polygons. Note that this option only works when you have only provided one input image. If multiple images are given (in WCS mode), then the full area covered by all the images has to be shown and the polygon excluded. This can lead to a very large area if large surveys like COSMOS are used. So Crop will abort and notify you. In such cases, it is best to crop out the larger region you want, then mask the smaller region with this option.

-s STR

Section of the input image which you want to be cropped. See Crop section syntax for a complete explanation on the syntax required for this input.


The column in a catalog to read as a coordinate. The value can be either the column number (starting from 1), or a match/search in the table meta-data, see Selecting table columns. This option must be called multiple times, depending on the number of dimensions in the input dataset. If it is called more than necessary, the extra columns (later calls to this option on the command-line or configuration files) will be ignored, see Configuration file precedence.


Column selection of crop file name. The value can be either the column number (starting from 1), or a match/search in the table meta-data, see Selecting table columns. This option can be used both in Image and WCS modes, and not a mandatory. When a column is given to this option, the final crop base file name will be taken from the contents of this column. The directory will be determined by the --output option (current directory if not given) and the value to --suffix will be appended. When this column isn’t given, the row number will be used instead.

Output options:


Square box width of region in the center of the image to check for blank values. If any of the pixels in this central region of a crop (defined by its center) are blank, then it will not be stored in an output file. If the value to this option is zero, no checking is done. This check is only applied when the cropped region(s) are defined by their center (not by the vertices, see Crop modes).

The units of the value are interpreted based on the --mode value (in WCS or pixel units). The ultimate checked region size (in pixels) will be an odd integer around the center (converted from WCS, or when an even number of pixels are given to this option). In WCS mode, the value can be given as fractions, for example if the WCS units are in degrees, 0.1/3600 will correspond to a check size of 0.1 arcseconds.

Because survey regions don’t often have a clean square or rectangle shape, some of the pixels on the sides of the survey FITS image don’t commonly have any data and are blank (see Blank pixels). So when the catalog was not generated from the input image, it often happens that the image does not have data over some of the points.

When the given center of a crop falls in such regions or outside the dataset, and this option has a non-zero value, no crop will be created. Therefore with this option, you can specify a width of a small box (3 pixels is often good enough) around the central pixel of the cropped image. You can check which crops were created and which weren’t from the command-line (if --quiet was not called, see Operating mode options), or in Crop’s log file (see Crop output).

-p STR

The suffix (or post-fix) of the output files for when you want all the cropped images to have a special ending. One case where this might be helpful is when besides the science images, you want the weight images (or exposure maps, which are also distributed with survey images) of the cropped regions too. So in one run, you can set the input images to the science images and --suffix=_s.fits. In the next run you can set the weight images as input and --suffix=_w.fits.


Pixels outside of the input image that are in the crop box will not be used. By default they are filled with blank values (depending on type), see Blank pixels. This option only applies only in Image mode, see Crop modes.


In float or double images, it is common to give the value of zero to blank pixels. If the input image type is one of these two types, such pixels will also be considered as blank. You can disable this behavior with this option, see Blank pixels.

Operating mode options:


Operate in Image mode or WCS mode when the input coordinates can be both image or WCS. The value must either be img or wcs, see Crop modes for a full description.

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The string given to --output option will be interpreted depending on how many crops were requested, see Crop modes:

The header of each output cropped image will contain the names of the input image(s) it was cut from. If a name is longer than the 70 character space that the FITS standard allows for header keyword values, the name will be cut into several keywords from the nearest slash (/). The keywords have the following format: ICFn_m (for Crop File). Where n is the number of the image used in this crop and m is the part of the name (it can be broken into multiple keywords). Following the name is another keyword named ICFnPIX which shows the pixel range from that input image in the same syntax as Crop section syntax. So this string can be directly given to the --section option later.

Once done, a log file can be created in the current directory with the --log option. This file will have three columns and the same number of rows as the number of cropped images. There are also comments on the top of the log file explaining basic information about the run and descriptions for the columns. A short description of the columns is also given below:

  1. The cropped image file name for that row.
  2. The number of input images that were used to create that image.
  3. A 0 if the central few pixels (value to the --checkcenter option) are blank and 1 if they aren’t. When the crop was not defined by its center (see Crop modes), or --checkcenter was given a value of 0 (see Invoking Crop), the center will not be checked and this column will be given a value of -1.

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6.2 Arithmetic

It is commonly necessary to do operations on some or all of the elements of a dataset independently (pixels in an image). For example, in the reduction of raw data it is necessary to subtract the Sky value (Sky value) from each image image. Later (once the images as warped into a single grid using Warp for example, see Warp), the images are co-added (the output pixel grid is the average of the pixels of the individual input images). Arithmetic is Gnuastro’s program for such operations on your datasets directly from the command-line. It currently uses the reverse polish or post-fix notation, see Reverse polish notation and will work on the native data types of the input images/data to reduce CPU and RAM resources, see Numeric data types. For more information on how to run Arithmetic, please see Invoking Arithmetic.

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6.2.1 Reverse polish notation

The most common notation for arithmetic operations is the infix notation where the operator goes between the two operands, for example \(4+5\). While the infix notation is the preferred way in most programming languages, currently Arithmetic does not use it since it will require parenthesis which can complicate the implementation of the code. In the near future we do plan to adopt this notation90, but for the time being (due to time constraints on the developers), Arithmetic uses the post-fix or reverse polish notation. The Wikipedia article provides some excellent explanation on this notation but here we will give a short summary here for self-sufficiency.

In the post-fix notation, the operator is placed after the operands, as we will see below this removes the need to define parenthesis for most ordinary operators. For example, instead of writing 5+6, we write 5 6 +. To easily understand how this notation works, you can think of each operand as a node in a first-in-first-out stack. Every time an operator is confronted, it pops the number of operands it needs from the top of the stack (so they don’t exist in the stack any more), does its operation and pushes the result back on top of the stack. So if you want the average of 5 and 6, you would write: 5 6 + 2 /. The operations that are done are:

  1. 5 is an operand, so it is pushed to the top of the stack.
  2. 6 is an operand, so it is pushed to the top of the stack.
  3. + is a binary operator, so pull the top two elements of the stack and perform addition on them (the order is \(5+6\) in the example above). The result is 11, push it on top of the stack.
  4. 2 is an operand so push it onto the top of the stack.
  5. / is a binary operator, so pull out the top two elements of the stack (top-most is 2, then 11) and divide the second one by the first.

In the Arithmetic program, the operands can be FITS images or numbers. As you can see, very complicated procedures can be created without the need for parenthesis or worrying about precedence. Even functions which take an arbitrary number of arguments can be defined in this notation. This is a very powerful notation and is used in languages like Postscript 91 (the programming language in Postscript and compiled into PDF files) uses this notation.

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6.2.2 Arithmetic operators

The recognized operators in Arithmetic are listed below. See Reverse polish notation for more on how the operators and operands should be ordered on the command-line. The operands to all operators can be a data array (for example a FITS image) or a number, the output will be an array or number according to the inputs. For example a number multiplied by an array will produce an array. The conditional operators will return pixel, or numerical values of 0 (false) or 1 (true) and stored in an unsigned char data type (see Numeric data types).


Addition, so “4 5 +” is equivalent to \(4+5\).


Subtraction, so “4 5 -” is equivalent to \(4-5\).


Multiplication, so “4 5 x” is equivalent to \(4\times5\).


Division, so “4 5 /” is equivalent to \(4/5\).


Modulo (remainder), so “3 2 %” is equivalent to \(1\). Note that the modulo operator only works on integer types.


Absolute value of first operand, so “4 abs” is equivalent to \(|4|\).


First operand to the power of the second, so “4.3 5f pow” is equivalent to \(4.3^{5}\). Currently pow will only work on single or double precision floating point numbers or images. To be sure that a number is read as a floating point (even if it doesn’t have any non-zero decimals) put an f after it.


The square root of the first operand, so “5 sqrt” is equivalent to \(\sqrt{5}\). The output type is determined from the input, so the output of this example will be 2 (since 5 doesn’t have any non-zero decimal digits). If you want 2.23607, run 5f sqrt instead, the f will ensure that a number will be read as a floating point number, even if it doesn’t have decimal digits. If the input image has an integer type, you should explicitly convert the image to floating point, for example a.fits float sqrt, see the type conversion operators below.


Natural logarithm of first operand, so “4 log” is equivalent to \(\ln(4)\). The output type is determined from the input, see the explanation under sqrt for more.


Base-10 logarithm of first operand, so “4 log10” is equivalent to \(\log(4)\). The output type is determined from the input, see the explanation under sqrt for more.


Minimum (non-blank) value in the top operand on the stack, so “a.fits minvalue” will push the the minimum pixel value in this image onto the stack. Therefore this operator is mainly intended for data (for example images), if the top operand is a number, this operator just returns it without any change. So note that when this operator acts on a single image, the output will no longer be an image, but a number. The output of this operand is in the same type as the input.


Maximum (non-blank) value of first operand in the same type, similar to minvalue.


Number of non-blank elements in first operand in the uint64 type, similar to minvalue.


Sum of non-blank elements in first operand in the float32 type, similar to minvalue.


Mean value of non-blank elements in first operand in the float32 type, similar to minvalue.


Standard deviation of non-blank elements in first operand in the float32 type, similar to minvalue.


Median of non-blank elements in first operand with the same type, similar to minvalue.


The first popped operand to this operator must be a positive integer number which specifies how many further operands should be popped from the stack. The given number of operands must have the same type and size. Each pixel of the output of this operator will be set to the minimum value of the given number of operands (images) in that pixel.

For example the following command will produce an image with the same size and type as the inputs but each output pixel is set to the minimum respective pixel value of the three input images.

$ astarithmetic a.fits b.fits c.fits 3 min

Important notes:


Similar to min, but the pixels of the output will contain the maximum of the respective pixels in all operands in the stack.


Similar to min, but the pixels of the output will contain the number of the respective non-blank pixels in all input operands.


Similar to min, but the pixels of the output will contain the sum of the respective pixels in all input operands.


Similar to min, but the pixels of the output will contain the mean (average) of the respective pixels in all operands in the stack.


Similar to min, but the pixels of the output will contain the standard deviation of the respective pixels in all operands in the stack.


Similar to min, but the pixels of the output will contain the median of the respective pixels in all operands in the stack.


Apply mean filtering (or moving average) on the input dataset. During mean filtering, each pixel (data element) is replaced by the mean value of all its surrounding pixels (excluding blank values). The number of surrounding pixels in each dimension (to calculate the mean) is determined through the earlier operands that have been pushed onto the stack prior to the input dataset. The number of necessary operands is determined by the dimensions of the input dataset (first popped operand). The order of the dimensions on the command-line is the order in FITS format. Here is one example:

$ astarithmetic 5 4 image.fits filter-mean

In this example, each pixel is replaced by the mean of a 5 by 4 box around it. The box is 5 pixels along the first FITS dimension (horizontal when viewed in ds9) and 4 pixels along the second FITS dimension (vertical).

Each pixel will be placed in the center of the box that the mean is calculated on. If the given width along a dimension is even, then the center is assumed to be between the pixels (not in the center of a pixel). When the pixel is close to the edge, the pixels of the box that fall outside the image are ignored. Therefore, on the edge, less points will be used in calculating the mean.

The final effect of mean filtering is to smooth the input image, it is essentially a convolution with a kernel that has identical values for all its pixels (is flat), see Convolution process.

Note that blank pixels will also be affected by this operator: if there are any non-blank elements in the box surrounding a blank pixel, in the filtered image, it will have the mean of the non-blank elements, therefore it won’t be blank any more. If blank elements are important for your analysis, you can use the isblank with the where operator to set them back to blank after filtering.


Apply median filtering on the input dataset. This is very similar to filter-mean, except that instead of the mean value of the box pixels, the median value is used to replace a pixel value. For more on how to use this operator, please see filter-mean.

The median is less susceptible to outliers compared to the mean. As a result, after median filtering, the pixel values will be more discontinuous than mean filtering.


Apply a \(\sigma\)-clipped mean filtering onto the input dataset. This is very similar to filter-mean, except that all outliers (identified by the \(\sigma\)-clipping algorithm) have been removed, see Sigma clipping for more on the basics of this algorithm. As described there, two extra input parameters are necessary for \(\sigma\)-clipping: the multiple of \(\sigma\) and the termination criteria. filter-sigclip-mean therefore needs to pop two other operands from the stack after the dimensions of the box.

For example the line below uses the same box size as the example of filter-mean. However, all elements in the box that are iteratively beyond \(3\sigma\) of the distribution’s median are removed from the final calculation of the mean until the change in \(\sigma\) is less than \(0.2\).

$ astarithmetic 3 0.2 5 4 image.fits filter-sigclip-mean

The median (which needs a sorted dataset) is necessary for \(\sigma\)-clipping, therefore filter-sigclip-mean can be significantly slower than filter-mean. However, if there are strong outliers in the dataset that you want to ignore (for example emission lines on a spectrum when finding the continuum), this is a much better solution.


Apply a \(\sigma\)-clipped median filtering onto the input dataset. This operator and its necessary operands are almost identical to filter-sigclip-mean, except that after \(\sigma\)-clipping, the median value (which is less affected by outliers than the mean) is added back to the stack.


Interpolate all the blank elements of the second popped operand with the median of its nearest non-blank neighbors. The number of the nearest non-blank neighbors used to calculate the median is given by the first popped operand. Note that the distance of the nearest non-blank neighbors is irrelevant in this interpolation.


Collapse the given dataset (second popped operand), by summing all elements along the first popped operand (a dimension in FITS standard: counting from one, from fastest dimension). The returned dataset has one dimension less compared to the input.

The output will have a double-precision floating point type irrespective of the input dataset’s type. Doing the operation in double-precision (64-bit) floating point will help the collapse (summation) be affected less by floating point errors. But afterwards, single-precision floating points are usually enough in real (noisy) datasets. So depending on the type of the input and its nature, it is recommended to use one of the type conversion operators on the returned dataset.

If any WCS is present, the returned dataset will also lack the respective dimension in its WCS matrix. Therefore, when the WCS is important for later processing, be sure that the input is aligned with the respective axises: all non-diagonal elements in the WCS matrix are zero.

One common application of this operator is the creation of pseudo broad-band or narrow-band 2D images from 3D data cubes. For example integral field unit (IFU) data products that have two spatial dimensions (first two FITS dimensions) and one spectral dimension (third FITS dimension). The command below will collapse the whole third dimension into a 2D array the size of the first two dimensions, and then convert the output to single-precision floating point (as discussed above).

$ astarithmetic cube.fits 3 collapse-sum float32

Similar to collapse-sum, but the returned dataset will be the mean value along the collapsed dimension, not the sum.


Similar to collapse-sum, but the returned dataset will be the number of non-blank values along the collapsed dimension. The output will have a 32-bit signed integer type. If the input dataset doesn’t have blank values, all the elements in the returned dataset will have a single value (the length of the collapsed dimension). Therefore this is mostly relevant when there are blank values in the dataset.


Similar to collapse-sum, but the returned dataset will have the same numeric type as the input and will contain the minimum value for each pixel along the collapsed dimension.


Similar to collapse-sum, but the returned dataset will have the same numeric type as the input and will contain the maximum value for each pixel along the collapsed dimension.


Erode the foreground pixels (with value 1) of the input dataset (second popped operand). The first popped operand is the connectivity (see description in connected-components). Erosion is simply a flipping of all foreground pixels (to background; with value 0) that are “touching” background pixels. “Touching” is defined by the connectivity. In effect, this carves off the outer borders of the foreground, making them thinner. This operator assumes a binary dataset (all pixels are 0 and 1).


Dilate the foreground pixels (with value 1) of the input dataset (second popped operand). The first popped operand is the connectivity (see description in connected-components). Erosion is simply a flipping of all background pixels (with value 0) to foreground that are “touching” foreground pixels. “Touching” is defined by the connectivity. In effect, this expands the outer borders of the foreground. This operator assumes a binary dataset (all pixels are 0 and 1).


Find the connected components in the input dataset (second popped operand). The first popped is the connectivity used in the connected components algorithm. The second popped operand is the dataset where connected components are to be found. It is assumed to be a binary image (with values of 0 or 1). It must have an 8-bit unsigned integer type which is the format produced by conditional operators. This operator will return a labeled dataset where the non-zero pixels in the input will be labeled with a counter (starting from 1).

The connectivity is a number between 1 and the number of dimensions in the dataset (inclusive). 1 corresponds to the weakest (symmetric) connectivity between elements and the number of dimensions the strongest. For example on a 2D image, a connectivity of 1 corresponds to 4-connected neighbors and 2 corresponds to 8-connected neighbors.

One example usage of this operator can be the identification of regions above a certain threshold, as in the command below. With this command, Arithmetic will first separate all pixels greater than 100 into a binary image (where pixels with a value of 1 are above that value). Afterwards, it will label all those that are connected.

$ astarithmetic in.fits 100 gt 2 connected-components

If your input dataset doesn’t have a binary type, but you know all its values are 0 or 1, you can use the uint8 operator (below) to convert it to binary.


Flip background (0) pixels surrounded by foreground (1) in a binary dataset. This operator takes two operands (similar to connected-components): the first popped operand is the connectivity (to define a hole) and the second is the binary (0 or 1 valued) dataset to fill holes in.


Invert an unsigned integer dataset. This is the only operator that ignores blank values (which are set to be the maximum values in the unsigned integer types).

This is useful in cases where the target(s) has(have) been imaged in absorption as raw formats (which are unsigned integer types). With this option, the maximum value for the given type will be subtracted from each pixel value, thus “inverting” the image, so the target(s) can be treated as emission. This can be useful when the higher-level analysis methods/tools only work on emission (positive skew in the noise, not negative).


Less than: If the second popped (or left operand in infix notation, see Reverse polish notation) value is smaller than the first popped operand, then this function will return a value of 1, otherwise it will return a value of 0. If both operands are images, then all the pixels will be compared with their counterparts in the other image. If only one operand is an image, then all the pixels will be compared with the the single value (number) of the other operand. Finally if both are numbers, then the output is also just one number (0 or 1). When the output is not a single number, it will be stored as an unsigned char type.


Less or equal: similar to lt (‘less than’ operator), but returning 1 when the second popped operand is smaller or equal to the first.


Greater than: similar to lt (‘less than’ operator), but returning 1 when the second popped operand is greater than the first.


Greater or equal: similar to lt (‘less than’ operator), but returning 1 when the second popped operand is larger or equal to the first.


Equality: similar to lt (‘less than’ operator), but returning 1 when the two popped operands are equal (to double precision floating point accuracy).


Non-Equality: similar to lt (‘less than’ operator), but returning 1 when the two popped operands are not equal (to double precision floating point accuracy).


Logical AND: returns 1 if both operands have a non-zero value and 0 if both are zero. Both operands have to be the same kind: either both images or both numbers.


Logical OR: returns 1 if either one of the operands is non-zero and 0 only when both operators are zero. Both operands have to be the same kind: either both images or both numbers.


Logical NOT: returns 1 when the operand is zero and 0 when the operand is non-zero. The operand can be an image or number, for an image, it is applied to each pixel separately.


Test for a blank value (see Blank pixels). In essence, this is very similar to the conditional operators: the output is either 1 or 0 (see the ‘less than’ operator above). The difference is that it only needs one operand. Because of the definition of a blank pixel, a blank value is not even equal to itself, so you cannot use the equal operator above to select blank pixels. See the “Blank pixels” box below for more on Blank pixels in Arithmetic.


Change the input (pixel) value where/if a certain condition holds. The conditional operators above can be used to define the condition. Three operands are required for where. The input format is demonstrated in this simplified example:

$ astarithmetic modify.fits binary.fits if-true.fits where

The value of any pixel in modify.fits that corresponds to a non-zero and non-blank pixel of binary.fits will be changed to the value of the same pixel in if-true.fits (this may also be a number). The 3rd and 2nd popped operands (modify.fits and binary.fits respectively, see Reverse polish notation) have to have the same dimensions/size. if-true.fits can be either a number, or have the same dimension/size as the other two.

The 2nd popped operand (binary.fits) has to have uint8 (or unsigned char in standard C) type (see Numeric data types). It is treated as a binary dataset (with only two values: zero and non-zero, hence the name binary.fits in this example). However, commonly you won’t be dealing with an actual FITS file of a condition/binary image. You will probably define the condition in the same run based on some other reference image and use the conditional and logical operators above to make a true/false (or one/zero) image for you internally. For example the case below:

$ astarithmetic in.fits reference.fits 100 gt new.fits where

In the example above, any of the in.fits pixels that has a value in reference.fits greater than 100, will be replaced with the corresponding pixel in new.fits. Effectively the reference.fits 100 gt part created the condition/binary image which was added to the stack (in memory) and later used by where. The command above is thus equivalent to these two commands:

$ astarithmetic reference.fits 100 gt --output=binary.fits
$ astarithmetic in.fits binary.fits new.fits where

Finally, the input operands are read and used independently, so you can use the same file more than once as any of the operands.

When the 1st popped operand to where (if-true.fits) is a single number, it may be a NaN value (or any blank value, depending on its type) like the example below (see Blank pixels). When the number is blank, it will be converted to the blank value of the type of the 3rd popped operand (in.fits). Hence, in the example below, all the pixels in reference.fits that have a value greater than 100, will become blank in the natural data type of in.fits (even though NaN values are only defined for floating point types).

$ astarithmetic in.fits reference.fits 100 gt nan where

Bitwise AND operator: only bits with values of 1 in both popped operands will get the value of 1, the rest will be set to 0. For example (assuming numbers can be written as bit strings on the command-line): 00101000 00100010 bitand will give 00100000. Note that the bitwise operators only work on integer type datasets.


Bitwise inclusive OR operator: The bits where at least one of the two popped operands has a 1 value get a value of 1, the others 0. For example (assuming numbers can be written as bit strings on the command-line): 00101000 00100010 bitand will give 00101010. Note that the bitwise operators only work on integer type datasets.


Bitwise exclusive OR operator: A bit will be 1 if it differs between the two popped operands. For example (assuming numbers can be written as bit strings on the command-line): 00101000 00100010 bitand will give 00001010. Note that the bitwise operators only work on integer type datasets.


Bitwise left shift operator: shift all the bits of the first operand to the left by a number of times given by the second operand. For example (assuming numbers can be written as bit strings on the command-line): 00101000 2 lshift will give 10100000. This is equivalent to multiplication by 4. Note that the bitwise operators only work on integer type datasets.


Bitwise right shift operator: shift all the bits of the first operand to the right by a number of times given by the second operand. For example (assuming numbers can be written as bit strings on the command-line): 00101000 2 rshift will give 00001010. Note that the bitwise operators only work on integer type datasets.


Bitwise not (more formally known as one’s complement) operator: flip all the bits of the popped operand (note that this is the only unary, or single operand, bitwise operator). In other words, any bit with a value of 0 is changed to 1 and vice-versa. For example (assuming numbers can be written as bit strings on the command-line): 00101000 bitnot will give 11010111. Note that the bitwise operators only work on integer type datasets/numbers.


Convert the type of the popped operand to 8-bit unsigned integer type (see Numeric data types). The internal conversion of C will be used.


Convert the type of the popped operand to 8-bit signed integer type (see Numeric data types). The internal conversion of C will be used.


Convert the type of the popped operand to 16-bit unsigned integer type (see Numeric data types). The internal conversion of C will be used.


Convert the type of the popped operand to 16-bit signed integer (see Numeric data types). The internal conversion of C will be used.


Convert the type of the popped operand to 32-bit unsigned integer type (see Numeric data types). The internal conversion of C will be used.


Convert the type of the popped operand to 32-bit signed integer type (see Numeric data types). The internal conversion of C will be used.


Convert the type of the popped operand to 64-bit unsigned integer (see Numeric data types). The internal conversion of C will be used.


Convert the type of the popped operand to 32-bit (single precision) floating point (see Numeric data types). The internal conversion of C will be used.


Convert the type of the popped operand to 64-bit (double precision) floating point (see Numeric data types). The internal conversion of C will be used.


Set the characters after the dash (AAA in the case shown here) as a name for the first popped operand on the stack. The named dataset will be freed from memory as soon as it is no longer needed, or if the name is reset to refer to another dataset later in the command. This operator thus enables re-usability of a dataset without having to re-read it from a file every time it is necessary during a process. When a dataset is necessary more than once, this operator can thus help simplify reading/writing on the command-line (thus avoiding potential bugs), while also speeding up the processing.

Like all operators, this operator pops the top operand off of the main processing stack, but unlike other operands, it won’t add anything back to the stack immediately. It will keep the popped dataset in memory through a separate list of named datasets (not on the main stack). That list will be used to add/copy any requested dataset to the main processing stack when the name is called.

The name to give the popped dataset is part of the operator’s name. For example the set-a operator of the command below, gives the name “a” to the contents of image.fits. This name is then used instead of the actual filename to multiply the dataset by two.

$ astarithmetic image.fits set-a a 2 x

The name can be any string, but avoid strings ending with standard filename suffixes (for example .fits)92.

One example of the usefulness of this operator is in the where operator. For example, let’s assume you want to mask all pixels larger than 5 in image.fits (extension number 1) with a NaN value. Without setting a name for the dataset, you have to read the file two times from memory in a command like this:

$ astarithmetic image.fits image.fits 5 gt nan where -g1

But with this operator you can simply give image.fits the name i and simplify the command above to the more readable one below (which greatly helps when the filename is long):

$ astarithmetic image.fits set-i   i i 5 gt nan where

Blank pixels in Arithmetic: Blank pixels in the image (see Blank pixels) will be stored based on the data type. When the input is floating point type, blank values are NaN. One aspect of NaN values is that by definition they will fail on any comparison. Hence both equal and not-equal operators will fail when both their operands are NaN! Therefore, the only way to guarantee selection of blank pixels is through the isblank operator explained above.

One way you can exploit this property of the NaN value to your advantage is when you want a fully zero-valued image (even over the blank pixels) based on an already existing image (with same size and world coordinate system settings). The following command will produce this for you:

$ astarithmetic input.fits nan eq --output=all-zeros.fits

Note that on the command-line you can write NaN in any case (for example NaN, or NAN are also acceptable). Reading NaN as a floating point number in Gnuastro isn’t case-sensitive.

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6.2.3 Invoking Arithmetic

Arithmetic will do pixel to pixel arithmetic operations on the individual pixels of input data and/or numbers. For the full list of operators with explanations, please see Arithmetic operators. Any operand that only has a single element (number, or single pixel FITS image) will be read as a number, the rest of the inputs must have the same dimensions. The general template is:

$ astarithmetic [OPTION...] ASTRdata1 [ASTRdata2] OPERATOR ...

One line examples:

## Calculate (10.32-3.84)^2.7 quietly (will just print 155.329):
$ astarithmetic -q 10.32 3.84 - 2.7 pow

## Inverse the input image (1/pixel):
$ astarithmetic 1 image.fits / --out=inverse.fits

## Multiply each pixel in image by -1:
$ astarithmetic image.fits -1 x --out=negative.fits

## Subtract extension 4 from extension 1 (counting from zero):
$ astarithmetic image.fits image.fits - --out=skysub.fits           \
                --hdu=1 --hdu=4

## Add two images, then divide them by 2 (2 is read as floating point):
$ astarithmetic image1.fits image2.fits + 2f / --out=average.fits

## Use Arithmetic's average operator:
$ astarithmetic image1.fits image2.fits average --out=average.fits

## Calculate the median of three images in three separate extensions:
$ astarithmetic img1.fits img2.fits img3.fits median                \
                -h0 -h1 -h2 --out=median.fits

If the output is an image, and the --output option is not given, automatic output will use the name of the first FITS image encountered to generate an output file name, see Automatic output. Also, output WCS information will be taken from the first input image encountered. When the output is a single number, that number will be printed in the standard output and no output file will be created. Arithmetic’s notation for giving operands to operators is described in Reverse polish notation. To ignore certain pixels, set them as blank, see Blank pixels, for example with the where operator (see Arithmetic operators). See Common options for a review of the options in all Gnuastro programs. Arithmetic just redefines the --hdu option as explained below:

--hdu INT/STR

The header data unit of the input FITS images, see Input/Output options. Unlike most options in Gnuastro (which will ultimately only have one value for this option), Arithmetic allows --hdu to be called multiple times and the value of each invocation will be stored separately (for the unlimited number of input images you would like to use). Recall that for other programs this (common) option only takes a single value. So in other programs, if you specify it multiple times on the command-line, only the last value will be used and in the configuration files, it will be ignored if it already has a value.

The order of the values to --hdu has to be in the same order as input FITS images. Options are first read from the command-line (from left to right), then top-down in each configuration file, see Configuration file precedence.

If the number of HDUs is less than the number of input images, Arithmetic will abort and notify you. However, if there are more HDUs than FITS images, there is no problem: they will be used in the given order (every time a FITS image comes up on the stack) and the extra HDUs will be ignored in the end. So there is no problem with having extra HDUs in the configuration files and by default several HDUs with a value of 0 are kept in the system-wide configuration file when you install Gnuastro.

--globalhdu INT/STR

Use the value to this option as the HDU of all input FITS files. This option is very convenient when you have many input files and the dataset of interest is in the same HDU of all the files. When this option is called, any values given to the --hdu option (explained above) are ignored and will not be used.

Arithmetic accepts two kinds of input: images and numbers. Images are considered to be any of the inputs that is a file name of a recognized type (see Arguments) and has more than one element/pixel. Numbers on the command-line will be read into the smallest type (see Numeric data types) that can store them, so -2 will be read as a char type (which is signed on most systems and can thus keep negative values), 2500 will be read as an unsigned short (all positive numbers will be read as unsigned), while 3.1415926535897 will be read as a double and 3.14 will be read as a float. To force a number to be read as float, add a f after it, so 5f will be added to the stack as float (see Reverse polish notation).

Unless otherwise stated (in Arithmetic operators), the operators can deal with numeric multiple data types (see Numeric data types). For example in “a.fits b.fits +”, the image types can be long and float. In such cases, C’s internal type conversion will be used. The output type will be set to the higher-ranking type of the two inputs. Unsigned integer types have smaller ranking than their signed counterparts and floating point types have higher ranking than the integer types. So the internal C type conversions done in the example above are equivalent to this piece of C:

size_t i;
long a[100];
float b[100], out[100];
for(i=0;i<100;++i) out[i]=a[i]+b[i];

Relying on the default C type conversion significantly speeds up the processing and also requires less RAM (when using very large images).

Some operators can only work on integer types (of any length, for example bitwise operators) while others only work on floating point types, (currently only the pow operator). In such cases, if the operand type(s) are different, an error will be printed. Arithmetic also comes with internal type conversion operators which you can use to convert the data into the appropriate type, see Arithmetic operators.

The hyphen (-) can be used both to specify options (see Options) and also to specify a negative number which might be necessary in your arithmetic. In order to enable you to do this, Arithmetic will first parse all the input strings and if the first character after a hyphen is a digit, then that hyphen is temporarily replaced by the vertical tab character which is not commonly used. The arguments are then parsed and these strings will not be specified as an option. Then the given arguments are parsed and any vertical tabs are replaced back with a hyphen so they can be read as negative numbers. Therefore, as long as the names of the files you want to work on, don’t start with a vertical tab followed by a digit, there is no problem. An important consequence of this implementation is that you should not write negative fractions like this: -.3, instead write them as -0.3.

Without any images, Arithmetic will act like a simple calculator and print the resulting output number on the standard output like the first example above. If you really want such calculator operations on the command-line, AWK (GNU AWK is the most common implementation) is much faster, easier and much more powerful. For example, the numerical one-line example above can be done with the following command. In general AWK is a fantastic tool and GNU AWK has a wonderful manual ( So if you often confront situations like this, or have to work with large text tables/catalogs, be sure to checkout AWK and simplify your life.

$ echo "" | awk '{print (10.32-3.84)^2.7}'

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6.3 Convolve

On an image, convolution can be thought of as a process to blur or remove the contrast in an image. If you are already familiar with the concept and just want to run Convolve, you can jump to Convolution kernel and Invoking Convolve and skip the lengthy introduction on the basic definitions and concepts of convolution.

There are generally two methods to convolve an image. The first and more intuitive one is in the “spatial domain” or using the actual image pixel values, see Spatial domain convolution. The second method is when we manipulate the “frequency domain”, or work on the magnitudes of the different frequencies that constitute the image, see Frequency domain and Fourier operations. Understanding convolution in the spatial domain is more intuitive and thus recommended if you are just starting to learn about convolution. However, getting a good grasp of the frequency domain is a little more involved and needs some concentration and some mathematical proofs. However, its reward is a faster operation and more importantly a very fundamental understanding of this very important operation.

Convolution of an image will generally result in blurring the image because it mixes pixel values. In other words, if the image has sharp differences in neighboring pixel values93, those sharp differences will become smoother. This has very good consequences in detection of signal in noise for example. In an actual observed image, the variation in neighboring pixel values due to noise can be very high. But after convolution, those variations will decrease and we have a better hope in detecting the possible underlying signal. Another case where convolution is extensively used is in mock images and modeling in general, convolution can be used to simulate the effect of the atmosphere or the optical system on the mock profiles that we create, see Point spread function. Convolution is a very interesting and important topic in any form of signal analysis (including astronomical observations). So we have thoroughly94 explained the concepts behind it in the following sub-sections.

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6.3.1 Spatial domain convolution

The pixels in an input image represent different “spatial” positions, therefore when convolution is done only using the actual input pixel values, we name the process as being done in the “Spatial domain”. In particular this is in contrast to the “frequency domain” that we will discuss later in Frequency domain and Fourier operations. In the spatial domain (and in realistic situations where the image and the convolution kernel don’t extend to infinity), convolution is the process of changing the value of one pixel to the weighted average of all the pixels in its neighborhood.

The ‘neighborhood’ of each pixel (how many pixels in which direction) and the ‘weight’ function (how much each neighboring pixel should contribute depending on its position) are given through a second image which is known as a “kernel”95.

Next: , Previous: , Up: Spatial domain convolution   [Contents][Index] Convolution process

In convolution, the kernel specifies the weight and positions of the neighbors of each pixel. To find the convolved value of a pixel, the central pixel of the kernel is placed on that pixel. The values of each overlapping pixel in the kernel and image are multiplied by each other and summed for all the kernel pixels. To have one pixel in the center, the sides of the convolution kernel have to be an odd number. This process effectively mixes the pixel values of each pixel with its neighbors, resulting in a blurred image compared to the sharper input image.

Formally, convolution is one kind of linear ‘spatial filtering’ in image processing texts. If we assume that the kernel has \(2a+1\) and \(2b+1\) pixels on each side, the convolved value of a pixel placed at \(x\) and \(y\) (\(C_{x,y}\)) can be calculated from the neighboring pixel values in the input image (\(I\)) and the kernel (\(K\)) from


Any pixel coordinate that is outside of the image in the equation above will be considered to be zero. When the kernel is symmetric about its center the blurred image has the same orientation as the original image. However, if the kernel is not symmetric, the image will be affected in the opposite manner, this is a natural consequence of the definition of spatial filtering. In order to avoid this we can rotate the kernel about its center by 180 degrees so the convolved output can have the same original orientation. Technically speaking, only if the kernel is flipped the process is known Convolution. If it isn’t it is known as Correlation.

To be a weighted average, the sum of the weights (the pixels in the kernel) have to be unity. This will have the consequence that the convolved image of an object and unconvolved object will have the same brightness (see Flux Brightness and magnitude), which is natural, because convolution should not eat up the object photons, it only disperses them.

Previous: , Up: Spatial domain convolution   [Contents][Index] Edges in the spatial domain

In purely ‘linear’ spatial filtering (convolution), there are problems on the edges of the input image. Here we will explain the problem in the spatial domain. For a discussion of this problem from the frequency domain perspective, see Edges in the frequency domain. The problem originates from the fact that on the edges, in practice96, the sum of the weights we use on the actual image pixels is not unity. For example, as discussed above, a profile in the center of an image will have the same brightness before and after convolution. However, for partially imaged profile on the edge of the image, the brightness (sum of its pixel fluxes within the image, see Flux Brightness and magnitude) will not be equal, some of the flux is going to be ‘eaten’ by the edges.

If you ran $ make check on the source files of Gnuastro, you can see the this effect by comparing the convolve_frequency.fits with convolve_spatial.fits in the ./tests/ directory. In the spatial domain, by default, no assumption will be made about pixels outside of the image or any blank pixels in the image. The problem explained above will also occur on the sides of blank regions (see Blank pixels). The solution to this edge effect problem is only possible in the spatial domain. For pixels near the edge, we have to abandon the assumption that the sum of the kernel pixels is unity during the convolution process97. So taking \(W\) as the sum of the kernel pixels that overlapped with non-blank and in-image pixels, the equation in Convolution process will become:

$$C_{x,y}= { \sum_{s=-a}^{a}\sum_{t=-b}^{b}K_{s,t}\times{}I_{x+s,y+t} \over W}.$$

In this manner, objects which are near the edges of the image or blank pixels will also have the same brightness (within the image) before and after convolution. This correction is applied by default in Convolve when convolving in the spatial domain. To disable it, you can use the --noedgecorrection option. In the frequency domain, there is no way to avoid this loss of flux near the edges of the image, see Edges in the frequency domain for an interpretation from the frequency domain perspective.

Note that the edge effect discussed here is different from the one in If convolving afterwards. In making mock images we want to simulate a real observation. In a real observation the images of the galaxies on the sides of the CCD are first blurred by the atmosphere and instrument, then imaged. So light from the parts of a galaxy which are immediately outside the CCD will affect the parts of the galaxy which are covered by the CCD. Therefore in modeling the observation, we have to convolve an image that is larger than the input image by exactly half of the convolution kernel. We can hence conclude that this correction for the edges is only useful when working on actual observed images (where we don’t have any more data on the edges) and not in modeling.

Next: , Previous: , Up: Convolve   [Contents][Index]

6.3.2 Frequency domain and Fourier operations

Getting a good grip on the frequency domain is usually not an easy job! So we have decided to give the issue a complete review here. Convolution in the frequency domain (see Convolution theorem) heavily relies on the concepts of Fourier transform (Fourier transform) and Fourier series (Fourier series) so we will be investigating these important operations first. It has become something of a cliché for people to say that the Fourier series “is a way to represent a (wave-like) function as the sum of simple sine waves” (from Wikipedia). However, sines themselves are abstract functions, so this statement really adds no extra layer of physical insight.

Before jumping head-first into the equations and proofs, we will begin with a historical background to see how the importance of frequencies actually roots in our ancient desire to see everything in terms of circles. A short review of how the complex plane should be interpreted is then given. Having paved the way with these two basics, we define the Fourier series and subsequently the Fourier transform. The final aim is to explain discrete Fourier transform, however some very important concepts need to be solidified first: The Dirac comb, convolution theorem and sampling theorem. So each of these topics are explained in their own separate sub-sub-section before going on to the discrete Fourier transform. Finally we revisit (after Edges in the spatial domain) the problem of convolution on the edges, but this time in the frequency domain. Understanding the sampling theorem and the discrete Fourier transform is very important in order to be able to pull out valuable science from the discrete image pixels. Therefore we have included the mathematical proofs and figures so you can have a clear understanding of these very important concepts.

Next: , Previous: , Up: Frequency domain and Fourier operations   [Contents][Index] Fourier series historical background

Ever since the ancient times, the circle has been (and still is) the simplest shape for abstract comprehension. All you need is a center point and a radius and you are done. All the points on a circle are at a fixed distance from the center. However, the moment you try to connect this elegantly simple and beautiful abstract construct (the circle) with the real world (for example compute its area or its circumference), things become really hard (ideally, impossible) because the irrational number \(\pi\) gets involved.

The key to understanding the Fourier series (thus the Fourier transform and finally the Discrete Fourier Transform) is our ancient desire to express everything in terms of circles or the most exceptionally simple and elegant abstract human construct. Most people prefer to say the same thing in a more ahistorical manner: to break a function into sines and cosines. As the term “ancient” in the previous sentence implies, Jean-Baptiste Joseph Fourier (1768 – 1830 A.D.) was not the first person to do this. The main reason we know this process by his name today is that he came up with an ingenious method to find the necessary coefficients (radius of) and frequencies (“speed” of rotation on) the circles for any generic (integrable) function.

Middle ages epicycles along
with two demonstrations of breaking a generic function using epicycles.

Figure 6.1: Epicycles and the Fourier series. Left: A demonstration of Mercury’s epicycles relative to the “center of the world” by Qutb al-Din al-Shirazi (1236 – 1311 A.D.) retrieved from Wikipedia. Middle and Right: How adding more epicycles (or terms in the Fourier series) will approximate functions. The right animation is also available.

Like most aspects of mathematics, this process of interpreting everything in terms of circles, began for astronomical purposes. When astronomers noticed that the orbit of Mars and other outer planets, did not appear to be a simple circle (as everything should have been in the heavens). At some point during their orbit, the revolution of these planets would become slower, stop, go back a little (in what is known as the retrograde motion) and then continue going forward again.

The correction proposed by Ptolemy (90 – 168 A.D.) was the most agreed upon. He put the planets on Epicycles or circles whose center itself rotates on a circle whose center is the earth. Eventually, as observations became more and more precise, it was necessary to add more and more epicycles in order to explain the complex motions of the planets98. Figure 6.1(Left) shows an example depiction of the epicycles of Mercury in the late 13th century.

Of course we now know that if they had abdicated the Earth from its throne in the center of the heavens and allowed the Sun to take its place, everything would become much simpler and true. But there wasn’t enough observational evidence for changing the “professional consensus” of the time to this radical view suggested by a small minority99. So the pre-Galilean astronomers chose to keep Earth in the center and find a correction to the models (while keeping the heavens a purely “circular” order).

The main reason we are giving this historical background which might appear off topic is to give historical evidence that while such “approximations” do work and are very useful for pragmatic reasons (like measuring the calendar from the movement of astronomical bodies). They offer no physical insight. The astronomers who were involved with the Ptolemaic world view had to add a huge number of epicycles during the centuries after Ptolemy in order to explain more accurate observations. Finally the death knell of this world-view was Galileo’s observations with his new instrument (the telescope). So the physical insight, which is what Astronomers and Physicists are interested in (as opposed to Mathematicians and Engineers who just like proving and optimizing or calculating!) comes from being creative and not limiting our selves to such approximations. Even when they work.

Next: , Previous: , Up: Frequency domain and Fourier operations   [Contents][Index] Circles and the complex plane

Before going onto the derivation, it is also useful to review how the complex numbers and their plane relate to the circles we talked about above. The two schematics in the middle and right of Figure 6.1 show how a 1D function of time can be made using the 2D real and imaginary surface. Seeing the animation in Wikipedia will really help in understanding this important concept. At each point in time, we take the vertical coordinate of the point and use it to find the value of the function at that point in time. Figure 6.2 shows this relation with the axes marked.

Leonhard Euler100 (1707 – 1783 A.D.) showed that the complex exponential (\(e^{iv}\) where \(v\) is real) is periodic and can be written as: \(e^{iv}=\cos{v}+isin{v}\). Therefore \(e^{iv+2\pi}=e^{iv}\). Later, Caspar Wessel (mathematician and cartographer 1745 – 1818 A.D.) showed how complex numbers can be displayed as vectors on a plane. Euler’s identity might seem counter intuitive at first, so we will try to explain it geometrically (for deeper physical insight). On the real-imaginary 2D plane (like the left hand plot in each box of Figure 6.2), multiplying a number by \(i\) can be interpreted as rotating the point by \(90\) degrees (for example the value \(3\) on the real axis becomes \(3i\) on the imaginary axis). On the other hand, \(e\equiv\lim_{n\rightarrow\infty}(1+{1\over n})^n\), therefore, defining \(m\equiv nu\), we get:

$$e^{u}=\lim_{n\rightarrow\infty}\left(1+{1\over n}\right)^{nu} =\lim_{n\rightarrow\infty}\left(1+{u\over nu}\right)^{nu} =\lim_{m\rightarrow\infty}\left(1+{u\over m}\right)^{m}$$

Taking \(u\equiv iv\) the result can be written as a generic complex number (a function of \(v\)):

$$e^{iv}=\lim_{m\rightarrow\infty}\left(1+i{v\over m}\right)^{m}=a(v)+ib(v)$$

For \(v=\pi\), a nice geometric animation of going to the limit can be seen on Wikipedia. We see that \(\lim_{m\rightarrow\infty}a(\pi)=-1\), while \(\lim_{m\rightarrow\infty}b(\pi)=0\), which gives the famous \(e^{i\pi}=-1\) equation. The final value is the real number \(-1\), however the distance of the polygon points traversed as \(m\rightarrow\infty\) is half the circumference of a circle or \(\pi\), showing how \(v\) in the equation above can be interpreted as an angle in units of radians and therefore how \(a(v)=cos(v)\) and \(b(v)=sin(v)\).

Since \(e^{iv}\) is periodic (let’s assume with a period of \(T\)), it is more clear to write it as \(v\equiv{2{\pi}n\over T}t\) (where \(n\) is an integer), so \(e^{iv}=e^{i{2{\pi}n\over T}t}\). The advantage of this notation is that the period (\(T\)) is clearly visible and the frequency (\(2{\pi}n \over T\), in units of 1/cycle) is defined through the integer \(n\). In this notation, \(t\) is in units of “cycle”s.

As we see from the examples in Figure 6.1 and Figure 6.2, for each constituting frequency, we need a respective ‘magnitude’ or the radius of the circle in order to accurately approximate the desired 1D function. The concepts of “period” and “frequency” are relatively easy to grasp when using temporal units like time because this is how we define them in every-day life. However, in an image (astronomical data), we are dealing with spatial units like distance. Therefore, by one “period” we mean the distance at which the signal is identical and frequency is defined as the inverse of that spatial “period”. The complex circle of Figure 6.2 can be thought of the Moon rotating about Earth which is rotating around the Sun; so the “Real (signal)” axis shows the Moon’s position as seen by a distant observer on the Sun as time goes by. Because of the scalar (not having any direction or vector) nature of time, Figure 6.2 is easier to understand in units of time. When thinking about spatial units, mentally replace the “Time (sec)” axis with “Distance (meters)”. Because length has direction and is a vector, visualizing the rotation of the imaginary circle and the advance along the “Distance (meters)” axis is not as simple as temporal units like time.


Figure 6.2: Relation between the real (signal), imaginary (\(i\equiv\sqrt{-1}\)) and time axes at two snapshots of time.

Next: , Previous: , Up: Frequency domain and Fourier operations   [Contents][Index] Fourier series

In astronomical images, our variable (brightness, or number of photo-electrons, or signal to be more generic) is recorded over the 2D spatial surface of a camera pixel. However to make things easier to understand, here we will assume that the signal is recorded in 1D (assume one row of the 2D image pixels). Also for this section and the next (Fourier transform) we will be talking about the signal before it is digitized or pixelated. Let’s assume that we have the continuous function \(f(l)\) which is integrable in the interval \([l_0, l_0+L]\) (always true in practical cases like images). Take \(l_0\) as the position of the first pixel in the assumed row of the image and \(L\) as the width of the image along that row. The units of \(l_0\) and \(L\) can be in any spatial units (for example meters) or an angular unit (like radians) multiplied by a fixed distance which is more common.

To approximate \(f(l)\) over this interval, we need to find a set of frequencies and their corresponding ‘magnitude’s (see Circles and the complex plane). Therefore our aim is to show \(f(l)\) as the following sum of periodic functions:

$$f(l)=\displaystyle\sum_{n=-\infty}^{\infty}c_ne^{i{2{\pi}n\over L}l} $$

Note that the different frequencies (\(2{\pi}n/L\), in units of cycles per meters for example) are not arbitrary. They are all integer multiples of the fundamental frequency of \(\omega_0=2\pi/L\). Recall that \(L\) was the length of the signal we want to model. Therefore, we see that the smallest possible frequency (or the frequency resolution) in the end, depends on the length we observed the signal or \(L\). In the case of each dimension on an image, this is the size of the image in the respective dimension. The frequencies have been defined in this “harmonic” fashion to insure that the final sum is periodic outside of the \([l_0, l_0+L]\) interval too. At this point, you might be thinking that the sky is not periodic with the same period as my camera’s view angle. You are absolutely right! The important thing is that since your camera’s observed region is the only region we are “observing” and will be using, the rest of the sky is irrelevant; so we can safely assume the sky is periodic outside of it. However, this working assumption will haunt us later in Edges in the frequency domain.

The frequencies are thus determined by definition. So all we need to do is to find the coefficients (\(c_n\)), or magnitudes, or radii of the circles for each frequency which is identified with the integer \(n\). Fourier’s approach was to multiply both sides with a fixed term:

$$f(l)e^{-i{2{\pi}m\over L}l}=\displaystyle\sum_{n=-\infty}^{\infty}c_ne^{i{2{\pi}(n-m)\over L}l} $$

where \(m>0\)101. We can then integrate both sides over the observation period:

$$\int_{l_0}^{l_0+L}f(l)e^{-i{2{\pi}m\over L}l}dl =\int_{l_0}^{l_0+L}\displaystyle\sum_{n=-\infty}^{\infty}c_ne^{i{2{\pi}(n-m)\over L}l}dl=\displaystyle\sum_{n=-\infty}^{\infty}c_n\int_{l_0}^{l_0+L}e^{i{2{\pi}(n-m)\over L}l}dl $$

Both \(n\) and \(m\) are positive integers. Also, we know that a complex exponential is periodic so after one period (\(L\)) it comes back to its starting point. Therefore \(\int_{l_0}^{l_0+L}e^{2{\pi}k/L}dl=0\) for any \(k>0\). However, when \(k=0\), this integral becomes: \(\int_{l_0}^{l_0+T}e^0dt=\int_{l_0}^{l_0+T}dt=T\). Hence since the integral will be zero for all \(n{\neq}m\), we get:

$$\displaystyle\sum_{n=-\infty}^{\infty}c_n\int_{l_0}^{l_0+T}e^{i{2{\pi}(n-m)\over L}l}dl=Lc_m $$

The origin of the axis is fundamentally an arbitrary position. So let’s set it to the start of the image such that \(l_0=0\). So we can find the “magnitude” of the frequency \(2{\pi}m/L\) within \(f(l)\) through the relation:

$$c_m={1\over L}\int_{0}^{L}f(l)e^{-i{2{\pi}m\over L}l}dl $$

Next: , Previous: , Up: Frequency domain and Fourier operations   [Contents][Index] Fourier transform

In Fourier series, we had to assume that the function is periodic outside of the desired interval with a period of \(L\). Therefore, assuming that \(L\rightarrow\infty\) will allow us to work with any function. However, with this approximation, the fundamental frequency (\(\omega_0\)) or the frequency resolution that we discussed in Fourier series will tend to zero: \(\omega_0\rightarrow0\). In the equation to find \(c_m\), every \(m\) represented a frequency (multiple of \(\omega_0\)) and the integration on \(l\) removes the dependence of the right side of the equation on \(l\), making it only a function of \(m\) or frequency. Let’s define the following two variables:

$$\omega{\equiv}m\omega_0={2{\pi}m\over L}$$


The equation to find the coefficients of each frequency in Fourier series thus becomes:

$$F(\omega)=\int_{-\infty}^{\infty}f(l)e^{-i{\omega}l}dl. $$

The function \(F(\omega)\) is thus the Fourier transform of \(f(l)\) in the frequency domain. So through this transformation, we can find (analyze) the magnitudes of the constituting frequencies or the value in the frequency space102 of our spatial input function. The great thing is that we can also do the reverse and later synthesize the input function from its Fourier transform. Let’s do it: with the approximations above, multiply the right side of the definition of the Fourier Series (Fourier series) with \(1=L/L=({\omega_0}L)/(2\pi)\):

$$f(l)={1\over 2\pi}\displaystyle\sum_{n=-\infty}^{\infty}Lc_ne^{{2{\pi}in\over L}l}\omega_0={1\over 2\pi}\displaystyle\sum_{n=-\infty}^{\infty}F(\omega)e^{i{\omega}l}\Delta\omega $$

To find the right most side of this equation, we renamed \(\omega_0\) as \(\Delta\omega\) because it was our resolution, \(2{\pi}n/L\) was written as \(\omega\) and finally, \(Lc_n\) was written as \(F(\omega)\) as we defined above. Now, as \(L\rightarrow\infty\), \(\Delta\omega\rightarrow0\) so we can write:

$$f(l)={1\over 2\pi}\int_{-\infty}^{\infty}F(\omega)e^{i{\omega}l}d\omega $$

Together, these two equations provide us with a very powerful set of tools that we can use to process (analyze) and recreate (synthesize) the input signal. Through the first equation, we can break up our input function into its constituent frequencies and analyze it, hence it is also known as analysis. Using the second equation, we can synthesize or make the input function from the known frequencies and their magnitudes. Thus it is known as synthesis. Here, we symbolize the Fourier transform (analysis) and its inverse (synthesis) of a function \(f(l)\) and its Fourier Transform \(F(\omega)\) as \({\cal F}[f]\) and \({\cal F}^{-1}[F]\).

Next: , Previous: , Up: Frequency domain and Fourier operations   [Contents][Index] Dirac delta and comb

The Dirac \(\delta\) (delta) function (also known as an impulse) is the way that we convert a continuous function into a discrete one. It is defined to satisfy the following integral:


When integrated with another function, it gives that function’s value at \(l=0\):


An impulse positioned at another point (say \(l_0\)) is written as \(\delta(l-l_0)\):


The Dirac \(\delta\) function also operates similarly if we use summations instead of integrals. The Fourier transform of the delta function is:

$${\cal F}[\delta(l)]=\int_{-\infty}^{\infty}\delta(l)e^{-i{\omega}l}dl=e^{-i{\omega}0}=1$$

$${\cal F}[\delta(l-l_0)]=\int_{-\infty}^{\infty}\delta(l-l_0)e^{-i{\omega}l}dl=e^{-i{\omega}l_0}$$

From the definition of the Dirac \(\delta\) we can also define a Dirac comb (\({\rm III}_P\)) or an impulse train with infinite impulses separated by \(P\):

$${\rm III}_P(l)\equiv\displaystyle\sum_{k=-\infty}^{\infty}\delta(l-kP) $$

\(P\) is chosen to represent “pixel width” later in Sampling theorem. Therefore the Dirac comb is periodic with a period of \(P\). We have intentionally used a different name for the period of the Dirac comb compared to the input signal’s length of observation that we showed with \(L\) in Fourier series. This difference is highlighted here to avoid confusion later when these two periods are needed together in Discrete Fourier transform. The Fourier transform of the Dirac comb will be necessary in Sampling theorem, so let’s derive it. By its definition, it is periodic, with a period of \(P\), so the Fourier coefficients of its Fourier Series (Fourier series) can be calculated within one period:

$${\rm III}_P=\displaystyle\sum_{n=-\infty}^{\infty}c_ne^{i{2{\pi}n\over P}l}$$

We can now find the \(c_n\) from Fourier series:

$$c_n={1\over P}\int_{-P/2}^{P/2}\delta(l)e^{-i{2{\pi}n\over P}l} ={1\over P}\quad\quad \rightarrow \quad\quad {\rm III}_P={1\over P}\displaystyle\sum_{n=-\infty}^{\infty}e^{i{2{\pi}n\over P}l} $$

So we can write the Fourier transform of the Dirac comb as:

$${\cal F}[{\rm III}_P]=\int_{-\infty}^{\infty}{\rm III}_Pe^{-i{\omega}l}dl ={1\over P}\displaystyle\sum_{n=-\infty}^{\infty}\int_{-\infty}^{\infty}e^{-i(\omega-{2{\pi}n\over P})l}dl={1\over P}\displaystyle\sum_{n=-\infty}^{\infty}\delta\left(\omega-{2{\pi}n\over P}\right) $$

In the last step, we used the fact that the complex exponential is a periodic function, that \(n\) is an integer and that as we defined in Fourier transform, \(\omega{\equiv}m\omega_0\), where \(m\) was an integer. The integral will be zero for any \(\omega\) that is not equal to \(2{\pi}n/P\), a more complete explanation can be seen in Fourier series. Therefore, while in the spatial domain the impulses had spacing of \(P\) (meters for example), in the frequency space, the spacing between the different impulses are \(2\pi/P\) cycles per meters.

Next: , Previous: , Up: Frequency domain and Fourier operations   [Contents][Index] Convolution theorem

The convolution (shown with the \(\ast\) operator) of the two functions \(f(l)\) and \(h(l)\) is defined as:

$$c(l)\equiv[f{\ast}h](l)=\int_{-\infty}^{\infty}f(\tau)h(l-\tau)d\tau $$

See Convolution process for a more detailed physical (pixel based) interpretation of this definition. The Fourier transform of convolution (\(C(\omega)\)) can be written as:

$$ C(\omega)=\int_{-\infty}^{\infty}[f{\ast}h](l)e^{-i{\omega}l}dl= \int_{-\infty}^{\infty}f(\tau)\left[\int_{-\infty}^{\infty}h(l-\tau)e^{-i{\omega}l}dl\right]d\tau $$

To solve the inner integral, let’s define \(s{\equiv}l-\tau\), so that \(ds=dl\) and \(l=s+\tau\) then the inner integral becomes:

$$\int_{-\infty}^{\infty}h(l-\tau)e^{-i{\omega}l}dl= \int_{-\infty}^{\infty}h(s)e^{-i{\omega}(s+\tau)}ds=e^{-i{\omega}\tau}\int_{-\infty}^{\infty}h(s)e^{-i{\omega}s}ds=H(\omega)e^{-i{\omega}\tau} $$

where \(H(\omega)\) is the Fourier transform of \(h(l)\). Substituting this result for the inner integral above, we get:

$$C(\omega)=H(\omega)\int_{-\infty}^{\infty}f(\tau)e^{-i{\omega}\tau}d\tau=H(\omega)F(\omega)=F(\omega)H(\omega) $$

where \(F(\omega)\) is the Fourier transform of \(f(l)\). So multiplying the Fourier transform of two functions individually, we get the Fourier transform of their convolution. The convolution theorem also proves a relation between the convolutions in the frequency space. Let’s define:


Applying the inverse Fourier Transform or synthesis equation (Fourier transform) to both sides and following the same steps above, we get:


Where \(d(l)\) is the inverse Fourier transform of \(D(\omega)\). We can therefore re-write the two equations above formally as the convolution theorem:

$$ {\cal F}[f{\ast}h]={\cal F}[f]{\cal F}[h] $$

$$ {\cal F}[fh]={\cal F}[f]\ast{\cal F}[h] $$

Besides its usefulness in blurring an image by convolving it with a given kernel, the convolution theorem also enables us to do another very useful operation in data analysis: to match the blur (or PSF) between two images taken with different telescopes/cameras or under different atmospheric conditions. This process is also known as de-convolution. Let’s take \(f(l)\) as the image with a narrower PSF (less blurry) and \(c(l)\) as the image with a wider PSF which appears more blurred. Also let’s take \(h(l)\) to represent the kernel that should be convolved with the sharper image to create the more blurry image. Above, we proved the relation between these three images through the convolution theorem. But there, we assumed that \(f(l)\) and \(h(l)\) are known (given) and the convolved image is desired.

In de-convolution, we have \(f(l)\) –the sharper image– and \(f*h(l)\) –the more blurry image– and we want to find the kernel \(h(l)\). The solution is a direct result of the convolution theorem:

$$ {\cal F}[h]={{\cal F}[f{\ast}h]\over {\cal F}[f]} \quad\quad {\rm or} \quad\quad h(l)={\cal F}^{-1}\left[{{\cal F}[f{\ast}h]\over {\cal F}[f]}\right] $$

While this works really nice, it has two problems:

A standard solution to both these problems is the Weiner de-convolution algorithm103.

Next: , Previous: , Up: Frequency domain and Fourier operations   [Contents][Index] Sampling theorem

Our mathematical functions are continuous, however, our data collecting and measuring tools are discrete. Here we want to give a mathematical formulation for digitizing the continuous mathematical functions so that later, we can retrieve the continuous function from the digitized recorded input. Assuming that we have a continuous function \(f(l)\), then we can define \(f_s(l)\) as the ‘sampled’ \(f(l)\) through the Dirac comb (see Dirac delta and comb):

$$f_s(l)=f(l){\rm III}_P=\displaystyle\sum_{n=-\infty}^{\infty}f(l)\delta(l-nP) $$

The discrete data-element \(f_k\) (for example, a pixel in an image), where \(k\) is an integer, can thus be represented as:


Note that in practice, our discrete data points are not found in this fashion. Each detector pixel (in an image for example) has an area and averages the signal it receives over that area, not a mathematical point as the Dirac \(\delta\) function defines. However, as long as the variation in the signal over one detector pixel is not significant, this can be a good approximation. Having put this issue to the side, we can now try to find the relation between the Fourier transforms of the unsampled \(f(l)\) and the sampled \(f_s(l)\). For a more clear notation, let’s define:

$$F_s(\omega)\equiv{\cal F}[f_s]$$

$$D(\omega)\equiv{\cal F}[{\rm III}_P]$$

Then using the Convolution theorem (see Convolution theorem), \(F_s(\omega)\) can be written as:

$$F_s(\omega)={\cal F}[f(l){\rm III}_P]=F(\omega){\ast}D(\omega)$$

Finally, from the definition of convolution and the Fourier transform of the Dirac comb (see Dirac delta and comb), we get:

$$\eqalign{ F_s(\omega) &= \int_{-\infty}^{\infty}F(\omega)D(\omega-\mu)d\mu \cr &= {1\over P}\displaystyle\sum_{n=-\infty}^{\infty}\int_{-\infty}^{\infty}F(\omega)\delta\left(\omega-\mu-{2{\pi}n\over P}\right)d\mu \cr &= {1\over P}\displaystyle\sum_{n=-\infty}^{\infty}F\left( \omega-{2{\pi}n\over P}\right).\cr } $$

\(F(\omega)\) was only a simple function, see Figure 6.3(left). However, from the sampled Fourier transform function we see that \(F_s(\omega)\) is the superposition of infinite copies of \(F(\omega)\) that have been shifted, see Figure 6.3(right). From the equation, it is clear that the shift in each copy is \(2\pi/P\).


Figure 6.3: Sampling causes infinite repetition in the frequency domain. FT is an abbreviation for ‘Fourier transform’. \(\omega_m\) represents the maximum frequency present in the input. \(F(\omega)\) is only symmetric on both sides of 0 when the input is real (not complex). In general \(F(\omega)\) is complex and thus cannot be simply plotted like this. Here we have assumed a real Gaussian \(f(t)\) which has produced a Gaussian \(F(\omega)\).

The input \(f(l)\) can have any distribution of frequencies in it. In the example of Figure 6.3(left), the input consisted of a range of frequencies equal to \(\Delta\omega=2\omega_m\). Fortunately as Figure 6.3(right) shows, the assumed pixel size (\(P\)) we used to sample this hypothetical function was such that \(2\pi/P>\Delta\omega\). The consequence is that each copy of \(F(\omega)\) has become completely separate from the surrounding copies. Such a digitized (sampled) data set is thus called over-sampled. When \(2\pi/P=\Delta\omega\), \(P\) is just small enough to finely separate even the largest frequencies in the input signal and thus it is known as critically-sampled. Finally if \(2\pi/P<\Delta\omega\) we are dealing with an under-sampled data set. In an under-sampled data set, the separate copies of \(F(\omega)\) are going to overlap and this will deprive us of recovering high constituent frequencies of \(f(l)\). The effects of under-sampling in an image with high rates of change (for example a brick wall imaged from a distance) can clearly be visually seen and is known as aliasing.

When the input \(f(l)\) is composed of a finite range of frequencies, \(f(l)\) is known as a band-limited function. The example in Figure 6.3(left) was a nice demonstration of such a case: for all \(\omega<-\omega_m\) or \(\omega>\omega_m\), we have \(F(\omega)=0\). Therefore, when the input function is band-limited and our detector’s pixels are placed such that we have critically (or over-) sampled it, then we can exactly reproduce the continuous \(f(l)\) from the discrete or digitized samples. To do that, we just have to isolate one copy of \(F(\omega)\) from the infinite copies and take its inverse Fourier transform.

This ability to exactly reproduce the continuous input from the sampled or digitized data leads us to the sampling theorem which connects the inherent property of the continuous signal (its maximum frequency) to that of the detector (the spacing between its pixels). The sampling theorem states that the full (continuous) signal can be recovered when the pixel size (\(P\)) and the maximum constituent frequency in the signal (\(\omega_m\)) have the following relation104:

$${2\pi\over P}>2\omega_m$$

This relation was first formulated by Harry Nyquist (1889 – 1976 A.D.) in 1928 and formally proved in 1949 by Claude E. Shannon (1916 – 2001 A.D.) in what is now known as the Nyquist-Shannon sampling theorem. In signal processing, the signal is produced (synthesized) by a transmitter and is received and de-coded (analyzed) by a receiver. Therefore producing a band-limited signal is necessary.

In astronomy, we do not produce the shapes of our targets, we are only observers. Galaxies can have any shape and size, therefore ideally, our signal is not band-limited. However, since we are always confined to observing through an aperture, the aperture will cause a point source (for which \(\omega_m=\infty\)) to be spread over several pixels. This spread is quantitatively known as the point spread function or PSF. This spread does blur the image which is undesirable; however, for this analysis it produces the positive outcome that there will be a finite \(\omega_m\). Though we should caution that any detector will have noise which will add lots of very high frequency (ideally infinite) changes between the pixels. However, the coefficients of those noise frequencies are usually exceedingly small.

Next: , Previous: , Up: Frequency domain and Fourier operations   [Contents][Index] Discrete Fourier transform

As we have stated several times so far, the input image is a digitized, pixelated or discrete array of values (\(f_s(l)\), see Sampling theorem). The input is not a continuous function. Also, all our numerical calculations can only be done on a sampled, or discrete Fourier transform. Note that \(F_s(\omega)\) is not discrete, it is continuous. One way would be to find the analytic \(F_s(\omega)\), then sample it at any desired “freq-pixel”105 spacing. However, this process would involve two steps of operations and computers in particular are not too good at analytic operations for the first step. So here, we will derive a method to directly find the ‘freq-pixel’ated \(F_s(\omega)\) from the pixelated \(f_s(l)\). Let’s start with the definition of the Fourier transform (see Fourier transform):

$$F_s(\omega)=\int_{-\infty}^{\infty}f_s(l)e^{-i{\omega}l}dl $$

From the definition of \(f_s(\omega)\) (using \(x\) instead of \(n\)) we get:

$$\eqalign{ F_s(\omega) &= \displaystyle\sum_{x=-\infty}^{\infty} \int_{-\infty}^{\infty}f(l)\delta(l-xP)e^{-i{\omega}l}dl \cr &= \displaystyle\sum_{x=-\infty}^{\infty} f_xe^{-i{\omega}xP} } $$

Where \(f_x\) is the value of \(f(l)\) on the point \(x\) or the value of the \(x\)th pixel. As shown in Sampling theorem this function is infinitely periodic with a period of \(2\pi/P\). So all we need is the values within one period: \(0<\omega<2\pi/P\), see Figure 6.3. We want \(X\) samples within this interval, so the frequency difference between each frequency sample or freq-pixel is \(1/XP\). Hence we will evaluate the equation above on the points at:

$$\omega={u\over XP} \quad\quad u = 0, 1, 2, ..., X-1$$

Therefore the value of the freq-pixel \(u\) in the frequency domain is:

$$F_u=\displaystyle\sum_{x=0}^{X-1} f_xe^{-i{ux\over X}} $$

Therefore, we see that for each freq-pixel in the frequency domain, we are going to need all the pixels in the spatial domain106. If the input (spatial) pixel row is also \(X\) pixels wide, then we can exactly recover the \(x\)th pixel with the following summation:

$$f_x={1\over X}\displaystyle\sum_{u=0}^{X-1} F_ue^{i{ux\over X}} $$

When the input pixel row (we are still only working on 1D data) has \(X\) pixels, then it is \(L=XP\) spatial units wide. \(L\), or the length of the input data was defined in Fourier series and \(P\) or the space between the pixels in the input was defined in Dirac delta and comb. As we saw in Sampling theorem, the input (spatial) pixel spacing (\(P\)) specifies the range of frequencies that can be studied and in Fourier series we saw that the length of the (spatial) input, (\(L\)) determines the resolution (or size of the freq-pixels) in our discrete Fourier transformed image. Both result from the fact that the frequency domain is the inverse of the spatial domain.

Next: , Previous: , Up: Frequency domain and Fourier operations   [Contents][Index] Fourier operations in two dimensions

Once all the relations in the previous sections have been clearly understood in one dimension, it is very easy to generalize them to two or even more dimensions since each dimension is by definition independent. Previously we defined \(l\) as the continuous variable in 1D and the inverse of the period in its direction to be \(\omega\). Let’s show the second spatial direction with \(m\) the the inverse of the period in the second dimension with \(\nu\). The Fourier transform in 2D (see Fourier transform) can be written as:

$$F(\omega, \nu)=\int_{-\infty}^{\infty}\int_{-\infty}^{\infty} f(l, m)e^{-i({\omega}l+{\nu}m)}dl$$

$$f(l, m)=\int_{-\infty}^{\infty}\int_{-\infty}^{\infty} F(\omega, \nu)e^{i({\omega}l+{\nu}m)}dl$$

The 2D Dirac \(\delta(l,m)\) is non-zero only when \(l=m=0\). The 2D Dirac comb (or Dirac brush! See Dirac delta and comb) can be written in units of the 2D Dirac \(\delta\). For most image detectors, the sides of a pixel are equal in both dimensions. So \(P\) remains unchanged, if a specific device is used which has non-square pixels, then for each dimension a different value should be used.

$${\rm III}_P(l, m)\equiv\displaystyle\sum_{j=-\infty}^{\infty} \displaystyle\sum_{k=-\infty}^{\infty} \delta(l-jP, m-kP) $$

The Two dimensional Sampling theorem (see Sampling theorem) is thus very easily derived as before since the frequencies in each dimension are independent. Let’s take \(\nu_m\) as the maximum frequency along the second dimension. Therefore the two dimensional sampling theorem says that a 2D band-limited function can be recovered when the following conditions hold107:

$${2\pi\over P} > 2\omega_m \quad\quad\quad {\rm and} \quad\quad\quad {2\pi\over P} > 2\nu_m$$

Finally, let’s represent the pixel counter on the second dimension in the spatial and frequency domains with \(y\) and \(v\) respectively. Also let’s assume that the input image has \(Y\) pixels on the second dimension. Then the two dimensional discrete Fourier transform and its inverse (see Discrete Fourier transform) can be written as:

$$F_{u,v}=\displaystyle\sum_{x=0}^{X-1}\displaystyle\sum_{y=0}^{Y-1} f_{x,y}e^{-i({ux\over X}+{vy\over Y})} $$

$$f_{x,y}={1\over XY}\displaystyle\sum_{u=0}^{X-1}\displaystyle\sum_{v=0}^{Y-1} F_{u,v}e^{i({ux\over X}+{vy\over Y})} $$

Previous: , Up: Frequency domain and Fourier operations   [Contents][Index] Edges in the frequency domain

With a good grasp of the frequency domain, we can revisit the problem of convolution on the image edges, see Edges in the spatial domain. When we apply the convolution theorem (see Convolution theorem) to convolve an image, we first take the discrete Fourier transforms (DFT, Discrete Fourier transform) of both the input image and the kernel, then we multiply them with each other and then take the inverse DFT to construct the convolved image. Of course, in order to multiply them with each other in the frequency domain, the two images have to be the same size, so let’s assume that we pad the kernel (it is usually smaller than the input image) with zero valued pixels in both dimensions so it becomes the same size as the input image before the DFT.

Having multiplied the two DFTs, we now apply the inverse DFT which is where the problem is usually created. If the DFT of the kernel only had values of 1 (unrealistic condition!) then there would be no problem and the inverse DFT of the multiplication would be identical with the input. However in real situations, the kernel’s DFT has a maximum of 1 (because the sum of the kernel has to be one, see Convolution process) and decreases something like the hypothetical profile of Figure 6.3. So when multiplied with the input image’s DFT, the coefficients or magnitudes (see Circles and the complex plane) of the smallest frequency (or the sum of the input image pixels) remains unchanged, while the magnitudes of the higher frequencies are significantly reduced.

As we saw in Sampling theorem, the Fourier transform of a discrete input will be infinitely repeated. In the final inverse DFT step, the input is in the frequency domain (the multiplied DFT of the input image and the kernel DFT). So the result (our output convolved image) will be infinitely repeated in the spatial domain. In order to accurately reconstruct the input image, we need all the frequencies with the correct magnitudes. However, when the magnitudes of higher frequencies are decreased, longer periods (shorter frequencies) will dominate in the reconstructed pixel values. Therefore, when constructing a pixel on the edge of the image, the newly empowered longer periods will look beyond the input image edges and will find the repeated input image there. So if you convolve an image in this fashion using the convolution theorem, when a bright object exists on one edge of the image, its blurred wings will be present on the other side of the convolved image. This is often termed as circular convolution or cyclic convolution.

So, as long as we are dealing with convolution in the frequency domain, there is nothing we can do about the image edges. The least we can do is to eliminate the ghosts of the other side of the image. So, we add zero valued pixels to both the input image and the kernel in both dimensions so the image that will be convolved has a size equal to the sum of both images in each dimension. Of course, the effect of this zero-padding is that the sides of the output convolved image will become dark. To put it another way, the edges are going to drain the flux from nearby objects. But at least it is consistent across all the edges of the image and is predictable. In Convolve, you can see the padded images when inspecting the frequency domain convolution steps with the --viewfreqsteps option.

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6.3.3 Spatial vs. Frequency domain

With the discussions above it might not be clear when to choose the spatial domain and when to choose the frequency domain. Here we will try to list the benefits of each.

The spatial domain,

The frequency domain,

As a general rule of thumb, when working on an image of modeled profiles use the frequency domain and when working on an image of real (observed) objects use the spatial domain (corrected for the edges). The reason is that if you apply a frequency domain convolution to a real image, you are going to loose information on the edges and generally you don’t want large kernels. But when you have made the profiles in the image yourself, you can just make a larger input image and crop the central parts to completely remove the edge effect, see If convolving afterwards. Also due to oversampling, both the kernels and the images can become very large and the speed boost of frequency domain convolution will significantly improve the processing time, see Oversampling.

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6.3.4 Convolution kernel

All the programs that need convolution will need to be given a convolution kernel file and extension. In most cases (other than Convolve, see Convolve) the kernel file name is optional. However, the extension is necessary and must be specified either on the command-line or at least one of the configuration files (see Configuration files). Within Gnuastro, there are two ways to create a kernel image:

The two options to specify a kernel file name and its extension are shown below. These are common between all the programs that will do convolution.

-k STR

The convolution kernel file name. The BITPIX (data type) value of this file can be any standard type and it does not necessarily have to be normalized. Several operations will be done on the kernel image prior to the program’s processing:


The convolution kernel HDU. Although the kernel file name is optional, before running any of the programs, they need to have a value for --khdu even if the default kernel is to be used. So be sure to keep its value in at least one of the configuration files (see Configuration files). By default, the system configuration file has a value.

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6.3.5 Invoking Convolve

Convolve an input image with a known kernel or make the kernel necessary to match two PSFs. The general template for Convolve is:

$ astconvolve [OPTION...] ASTRdata

One line examples:

## Convolve mockimg.fits with psf.fits:
$ astconvolve --kernel=psf.fits mockimg.fits

## Convolve in the spatial domain:
$ astconvolve observedimg.fits --kernel=psf.fits --domain=spatial

## Find the kernel to match sharper and blurry PSF images:
$ astconvolve --kernel=sharperimage.fits --makekernel=10           \

The only argument accepted by Convolve is an input image file. Some of the options are the same between Convolve and some other Gnuastro programs. Therefore, to avoid repetition, they will not be repeated here. For the full list of options shared by all Gnuastro programs, please see Common options. In particular, in the spatial domain convolve uses Gnuastro’s tessellation, see Tessellation and the common options related to that in Processing options.

Here we will only explain the options particular to Convolve. Run Convolve with --help in order to see the full list of options Convolve accepts, irrespective of where they are explained in this book.


Do not flip the kernel after reading it the spatial domain convolution. This can be useful if the flipping has already been applied to the kernel.


Do not normalize the kernel after reading it, such that the sum of its pixels is unity.

-d STR

The domain to use for the convolution. The acceptable values are ‘spatial’ and ‘frequency’, corresponding to the respective domain.

For large images, the frequency domain process will be more efficient than convolving in the spatial domain. However, the edges of the image will loose some flux (see Edges in the spatial domain) and the image must not contain any blank pixels, see Spatial vs. Frequency domain.


With this option a file with the initial name of the output file will be created that is suffixed with _freqsteps.fits, all the steps done to arrive at the final convolved image are saved as extensions in this file. The extensions in order are:

  1. The padded input image. In frequency domain convolution the two images (input and convolved) have to be the same size and both should be padded by zeros.
  2. The padded kernel, similar to the above.
  3. The Fourier spectrum of the forward Fourier transform of the input image. Note that the Fourier transform is a complex operation (and not view able in one image!) So we either have to show the ‘Fourier spectrum’ or the ‘Phase angle’. For the complex number \(a+ib\), the Fourier spectrum is defined as \(\sqrt{a^2+b^2}\) while the phase angle is defined as \(\arctan(b/a)\).
  4. The Fourier spectrum of the forward Fourier transform of the kernel image.
  5. The Fourier spectrum of the multiplied (through complex arithmetic) transformed images.
  6. The inverse Fourier transform of the multiplied image. If you open it, you will see that the convolved image is now in the center, not on one side of the image as it started with (in the padded image of the first extension). If you are working on a mock image which originally had pixels of precisely 0.0, you will notice that in those parts that your convolved profile(s) did not convert, the values are now \(\sim10^{-18}\), this is due to floating-point round off errors. Therefore in the final step (when cropping the central parts of the image), we also remove any pixel with a value less than \(10^{-17}\).

Do not correct the edge effect in spatial domain convolution. For a full discussion, please see Edges in the spatial domain.

-m INT

(=INT) If this option is called, Convolve will do de-convolution (see Convolution theorem). The image specified by the --kernel option is assumed to be the sharper (less blurry) image and the input image is assumed to be the more blurry image. The value given to this option will be used as the maximum radius of the kernel. Any pixel in the final kernel that is larger than this distance from the center will be set to zero. The two images must have the same size.

Noise has large frequencies which can make the result less reliable for the higher frequencies of the final result. So all the frequencies which have a spectrum smaller than the value given to the minsharpspec option in the sharper input image are set to zero and not divided. This will cause the wings of the final kernel to be flatter than they would ideally be which will make the convolved image result unreliable if it is too high. Some notes to take into account for a good result:


(=FLT) The minimum frequency spectrum (or coefficient, or pixel value in the frequency domain image) to use in deconvolution, see the explanations under the --makekernel option for more information.

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6.4 Warp

Image warping is the process of mapping the pixels of one image onto a new pixel grid. This process is sometimes known as transformation, however following the discussion of Heckbert 1989108 we will not be using that term because it can be confused with only pixel value or flux transformations. Here we specifically mean the pixel grid transformation which is better conveyed with ‘warp’.

Image wrapping is a very important step in astronomy, both in observational data analysis and in simulating modeled images. In modeling, warping an image is necessary when we want to apply grid transformations to the initial models, for example in simulating gravitational lensing (Radial warpings are not yet included in Warp). Observational reasons for warping an image are listed below:

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6.4.1 Warping basics

Let’s take \(\left[\matrix{u&v}\right]\) as the coordinates of a point in the input image and \(\left[\matrix{x&y}\right]\) as the coordinates of that same point in the output image110. The simplest form of coordinate transformation (or warping) is the scaling of the coordinates, let’s assume we want to scale the first axis by \(M\) and the second by \(N\), the output coordinates of that point can be calculated by

$$\left[\matrix{x\cr y}\right]= \left[\matrix{Mu\cr Nv}\right]= \left[\matrix{M&0\cr0&N}\right]\left[\matrix{u\cr v}\right]$$

Note that these are matrix multiplications. We thus see that we can represent any such grid warping as a matrix. Another thing we can do with this \(2\times2\) matrix is to rotate the output coordinate around the common center of both coordinates. If the output is rotated anticlockwise by \(\theta\) degrees from the positive (to the right) horizontal axis, then the warping matrix should become:

$$\left[\matrix{x\cr y}\right]= \left[\matrix{ucos\theta-vsin\theta\cr usin\theta+vcos\theta}\right]= \left[\matrix{cos\theta&-sin\theta\cr sin\theta&cos\theta}\right] \left[\matrix{u\cr v}\right] $$

We can also flip the coordinates around the first axis, the second axis and the coordinate center with the following three matrices respectively:

$$\left[\matrix{1&0\cr0&-1}\right]\quad\quad \left[\matrix{-1&0\cr0&1}\right]\quad\quad \left[\matrix{-1&0\cr0&-1}\right]$$

The final thing we can do with this definition of a \(2\times2\) warping matrix is shear. If we want the output to be sheared along the first axis with \(A\) and along the second with \(B\), then we can use the matrix:

$$\left[\matrix{1&A\cr B&1}\right]$$

To have one matrix representing any combination of these steps, you use matrix multiplication, see Merging multiple warpings. So any combinations of these transformations can be displayed with one \(2\times2\) matrix:

$$\left[\matrix{a&b\cr c&d}\right]$$

The transformations above can cover a lot of the needs of most coordinate transformations. However they are limited to mapping the point \([\matrix{0&0}]\) to \([\matrix{0&0}]\). Therefore they are useless if you want one coordinate to be shifted compared to the other one. They are also space invariant, meaning that all the coordinates in the image will receive the same transformation. In other words, all the pixels in the output image will have the same area if placed over the input image. So transformations which require varying output pixel sizes like projections cannot be applied through this \(2\times2\) matrix either (for example for the tilted ACS and WFC3 camera detectors on board the Hubble space telescope).

To add these further capabilities, namely translation and projection, we use the homogeneous coordinates. They were defined about 200 years ago by August Ferdinand Möbius (1790 – 1868). For simplicity, we will only discuss points on a 2D plane and avoid the complexities of higher dimensions. We cannot provide a deep mathematical introduction here, interested readers can get a more detailed explanation from Wikipedia111 and the references therein.

By adding an extra coordinate to a point we can add the flexibility we need. The point \([\matrix{x&y}]\) can be represented as \([\matrix{xZ&yZ&Z}]\) in homogeneous coordinates. Therefore multiplying all the coordinates of a point in the homogeneous coordinates with a constant will give the same point. Put another way, the point \([\matrix{x&y&Z}]\) corresponds to the point \([\matrix{x/Z&y/Z}]\) on the constant \(Z\) plane. Setting \(Z=1\), we get the input image plane, so \([\matrix{u&v&1}]\) corresponds to \([\matrix{u&v}]\). With this definition, the transformations above can be generally written as:

$$\left[\matrix{x\cr y\cr 1}\right]= \left[\matrix{a&b&0\cr c&d&0\cr 0&0&1}\right] \left[\matrix{u\cr v\cr 1}\right]$$

We thus acquired 4 extra degrees of freedom. By giving non-zero values to the zero valued elements of the last column we can have translation (try the matrix multiplication!). In general, any coordinate transformation that is represented by the matrix below is known as an affine transformation112:

$$\left[\matrix{a&b&c\cr d&e&f\cr 0&0&1}\right]$$

We can now consider translation, but the affine transform is still spatially invariant. Giving non-zero values to the other two elements in the matrix above gives us the projective transformation or Homography113 which is the most general type of transformation with the \(3\times3\) matrix:

$$\left[\matrix{x'\cr y'\cr w}\right]= \left[\matrix{a&b&c\cr d&e&f\cr g&h&1}\right] \left[\matrix{u\cr v\cr 1}\right]$$

So the output coordinates can be calculated from:

$$x={x' \over w}={au+bv+c \over gu+hv+1}\quad\quad\quad\quad y={y' \over w}={du+ev+f \over gu+hv+1}$$

Thus with Homography we can change the sizes of the output pixels on the input plane, giving a ‘perspective’-like visual impression. This can be quantitatively seen in the two equations above. When \(g=h=0\), the denominator is independent of \(u\) or \(v\) and thus we have spatial invariance. Homography preserves lines at all orientations. A very useful fact about Homography is that its inverse is also a Homography. These two properties play a very important role in the implementation of this transformation. A short but instructive and illustrated review of affine, projective and also bi-linear mappings is provided in Heckbert 1989114.

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6.4.2 Merging multiple warpings

In Warping basics we saw how a basic warp/transformation can be represented with a matrix. To make more complex warpings (for example to define a translation, rotation and scale as one warp) the individual matrices have to be multiplied through matrix multiplication. However matrix multiplication is not commutative, so the order of the set of matrices you use for the multiplication is going to be very important.

The first warping should be placed as the left-most matrix. The second warping to the right of that and so on. The second transformation is going to occur on the warped coordinates of the first. As an example for merging a few transforms into one matrix, the multiplication below represents the rotation of an image about a point \([\matrix{U&V}]\) anticlockwise from the horizontal axis by an angle of \(\theta\). To do this, first we take the origin to \([\matrix{U&V}]\) through translation. Then we rotate the image, then we translate it back to where it was initially. These three operations can be merged in one operation by calculating the matrix multiplication below:

$$\left[\matrix{1&0&U\cr0&1&V\cr{}0&0&1}\right] \left[\matrix{cos\theta&-sin\theta&0\cr sin\theta&cos\theta&0\cr 0&0&1}\right] \left[\matrix{1&0&-U\cr0&1&-V\cr{}0&0&1}\right]$$

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6.4.3 Resampling

A digital image is composed of discrete ‘picture elements’ or ‘pixels’. When a real image is created from a camera or detector, each pixel’s area is used to store the number of photo-electrons that were created when incident photons collided with that pixel’s surface area. This process is called the ‘sampling’ of a continuous or analog data into digital data. When we change the pixel grid of an image or warp it as we defined in Warping basics, we have to ‘guess’ the flux value of each pixel on the new grid based on the old grid, or re-sample it. Because of the ‘guessing’, any form of warping on the data is going to degrade the image and mix the original pixel values with each other. So if an analysis can be done on an unwarped data image, it is best to leave the image untouched and pursue the analysis. However as discussed in Warp this is not possible most of the times, so we have to accept the problem and re-sample the image.

In most applications of image processing, it is sufficient to consider each pixel to be a point and not an area. This assumption can significantly speed up the processing of an image and also the simplicity of the code. It is a fine assumption when the signal to noise ratio of the objects are very large. The question will then be one of interpolation because you have multiple points distributed over the output image and you want to find the values at the pixel centers. To increase the accuracy, you might also sample more than one point from within a pixel giving you more points for a more accurate interpolation in the output grid.

However, interpolation has several problems. The first one is that it will depend on the type of function you want to assume for the interpolation. For example you can choose a bi-linear or bi-cubic (the ‘bi’s are for the 2 dimensional nature of the data) interpolation method. For the latter there are various ways to set the constants115. Such functional interpolation functions can fail seriously on the edges of an image. They will also need normalization so that the flux of the objects before and after the warpings are comparable. The most basic problem with such techniques is that they are based on a point while a detector pixel is an area. They add a level of subjectivity to the data (make more assumptions through the functions than the data can handle). For most applications this is fine, but in scientific applications where detection of the faintest possible galaxies or fainter parts of bright galaxies is our aim, we cannot afford this loss. Because of these reasons Warp will not use such interpolation techniques.

Warp will do interpolation based on “pixel mixing”116 or “area resampling”. This is also what the Hubble Space Telescope pipeline calls “Drizzling”117. This technique requires no functions, it is thus non-parametric. It is also the closest we can get (make least assumptions) to what actually happens on the detector pixels. The basic idea is that you reverse-transform each output pixel to find which pixels of the input image it covers and what fraction of the area of the input pixels are covered. To find the output pixel value, you simply sum the value of each input pixel weighted by the overlap fraction (between 0 to 1) of the output pixel and that input pixel. Through this process, pixels are treated as an area not as a point (which is how detectors create the image), also the brightness (see Flux Brightness and magnitude) of an object will be left completely unchanged.

If there are very high spatial-frequency signals in the image (for example fringes) which vary on a scale smaller than your output image pixel size, pixel mixing can cause ailiasing118. So if the input image has fringes, they have to be calculated and removed separately (which would naturally be done in any astronomical application). Because of the PSF no astronomical target has a sharp change in the signal so this issue is less important for astronomical applications, see Point spread function.

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6.4.4 Invoking Warp

Warp an input dataset into a new grid. Any homographic warp (for example scaling, rotation, translation, projection) is acceptable, see Warping basics for the definitions. The general template for invoking Warp is:

$ astwarp [OPTIONS...] InputImage

One line examples:

## Rotate and then scale input image:
$ astwarp --rotate=37.92 --scale=0.8 image.fits

## Scale, then translate the input image:
$ astwarp --scale 8/3 --translate 2.1 image.fits

## Align raw image with celestial coordinates:
$ astwarp --align rawimage.fits --output=aligned.fits

## Directly input a custom warping matrix (using fraction):
$ astwarp --matrix=1/5,0,4/10,0,1/5,4/10,0,0,1 image.fits

## Directly input a custom warping matrix, with final numbers:
$ astwarp --matrix="0.7071,-0.7071,  0.7071,0.7071" image.fits

If any processing is to be done, Warp can accept one file as input. As in all Gnuastro programs, when an output is not explicitly set with the --output option, the output filename will be set automatically based on the operation, see Automatic output. For the full list of general options to all Gnuastro programs (including Warp), please see Common options.

To be the most accurate, the input image will be read as a 64-bit double precision floating point dataset and all internal processing is done in this format (including the raw output type). You can use the common --type option to write the output in any type you want, see Numeric data types.

Warps must be specified as command-line options, either as (possibly multiple) modular warpings (for example --rotate, or --scale), or directly as a single raw matrix (with --matrix). If specified together, the latter (direct matrix) will take precedence and all the modular warpings will be ignored. Any number of modular warpings can be specified on the command-line and configuration files. If more than one modular warping is given, all will be merged to create one warping matrix. As described in Merging multiple warpings, matrix multiplication is not commutative, so the order of specifying the modular warpings on the command-line, and/or configuration files makes a difference (see Configuration file precedence). The full list of modular warpings and the other options particular to Warp are described below.

The values to the warping options (modular warpings as well as --matrix), are a sequence of at least one number. Each number in this sequence is separated from the next by a comma (,). Each number can also be written as a single fraction (with a forward-slash / between the numerator and denominator). Space and Tab characters are permitted between any two numbers, just don’t forget to quote the whole value. Otherwise, the value will not be fully passed onto the option. See the examples above as a demonstration.

Based on the FITS standard, integer values are assigned to the center of a pixel and the coordinate [1.0, 1.0] is the center of the first pixel (bottom left of the image when viewed in SAO ds9). So the coordinate center [0.0, 0.0] is half a pixel away (in each axis) from the bottom left vertex of the first pixel. The resampling that is done in Warp (see Resampling) is done on the coordinate axes and thus directly depends on the coordinate center. In some situations this if fine, for example when rotating/aligning a real image, all the edge pixels will be similarly affected. But in other situations (for example when scaling an over-sampled mock image to its intended resolution, this is not desired: you want the center of the coordinates to be on the corner of the pixel. In such cases, you can use the --centeroncorner option which will shift the center by \(0.5\) before the main warp, then shift it back by \(-0.5\) after the main warp, see below.


Align the image and celestial (WCS) axes given in the input. After it, the vertical image direction (when viewed in SAO ds9) corresponds to the declination and the horizontal axis is the inverse of the Right Ascension (RA). The inverse of the RA is chosen so the image can correspond to what you would actually see on the sky and is common in most survey images.

Align is internally treated just like a rotation (--rotation), but uses the input image’s WCS to find the rotation angle. Thus, if you have rotated the image before calling --align, you might get unexpected results (because the rotation is defined on the original WCS).

-r FLT

Rotate the input image by the given angle in degrees: \(\theta\) in Warping basics. Note that commonly, the WCS structure of the image is set such that the RA is the inverse of the image horizontal axis which increases towards the right in the FITS standard and as viewed by SAO ds9. So the default center for rotation is on the right of the image. If you want to rotate about other points, you have to translate the warping center first (with --translate) then apply your rotation and then return the center back to the original position (with another call to --translate, see Merging multiple warpings.

-s FLT[,FLT]

Scale the input image by the given factor(s): \(M\) and \(N\) in Warping basics. If only one value is given, then both image axes will be scaled with the given value. When two values are given (separated by a comma), the first will be used to scale the first axis and the second will be used for the second axis. If you only need to scale one axis, use 1 for the axis you don’t need to scale. The value(s) can also be written (on the command-line or in configuration files) as a fraction.

-f FLT[,FLT]

Flip the input image around the given axis(s). If only one value is given, then both image axes are flipped. When two values are given (separated by a comma), you can choose which axis to flip over. --flip only takes values 0 (for no flip), or 1 (for a flip). Hence, if you want to flip by the second axis only, use --flip=0,1.

-e FLT[,FLT]

Shear the input image by the given value(s): \(A\) and \(B\) in Warping basics. If only one value is given, then both image axes will be sheared with the given value. When two values are given (separated by a comma), the first will be used to shear the first axis and the second will be used for the second axis. If you only need to shear along one axis, use 0 for the axis that must be untouched. The value(s) can also be written (on the command-line or in configuration files) as a fraction.

-t FLT[,FLT]

Translate (move the center of coordinates) the input image by the given value(s): \(c\) and \(f\) in Warping basics. If only one value is given, then both image axes will be translated by the given value. When two values are given (separated by a comma), the first will be used to translate the first axis and the second will be used for the second axis. If you only need to translate along one axis, use 0 for the axis that must be untouched. The value(s) can also be written (on the command-line or in configuration files) as a fraction.

-p FLT[,FLT]

Apply a projection to the input image by the given values(s): \(g\) and \(h\) in Warping basics. If only one value is given, then projection will apply to both axes with the given value. When two values are given (separated by a comma), the first will be used to project the first axis and the second will be used for the second axis. If you only need to project along one axis, use 0 for the axis that must be untouched. The value(s) can also be written (on the command-line or in configuration files) as a fraction.

-m STR

The warp/transformation matrix. All the elements in this matrix must be separated by comas(,) characters and as described above, you can also use fractions (a forward-slash between two numbers). The transformation matrix can be either a 2 by 2 (4 numbers), or a 3 by 3 (9 numbers) array. In the former case (if a 2 by 2 matrix is given), then it is put into a 3 by 3 matrix (see Warping basics).

The determinant of the matrix has to be non-zero and it must not contain any non-number values (for example infinities or NaNs). The elements of the matrix have to be written row by row. So for the general Homography matrix of Warping basics, it should be called with --matrix=a,b,c,d,e,f,g,h,1.

The raw matrix takes precedence over all the modular warping options listed above, so if it is called with any number of modular warps, the latter are ignored.


Put the center of coordinates on the corner of the first (bottom-left when viewed in SAO ds9) pixel. This option is applied after the final warping matrix has been finalized: either through modular warpings or the raw matrix. See the explanation above for coordinates in the FITS standard to better understand this option and when it should be used.


Specify the first header keyword number (line) that should be used to read the WCS information, see the full explanation in Invoking Crop.


Specify the last header keyword number (line) that should be used to read the WCS information, see the full explanation in Invoking Crop.


Do not correct the WCS information of the input image and save it untouched to the output image. By default the WCS (World Coordinate System) information of the input image is going to be corrected in the output image so the objects in the image are at the same WCS coordinates. But in some cases it might be useful to keep it unchanged (for example to correct alignments).


Depending on the warp, the output pixels that cover pixels on the edge of the input image, or blank pixels in the input image, are not going to be fully covered by input data. With this option, you can specify the acceptable covered fraction of such pixels (any value between 0 and 1). If you only want output pixels that are fully covered by the input image area (and are not blank), then you can set --coveredfrac=1. Alternatively, a value of 0 will keep output pixels that are even infinitesimally covered by the input(so the sum of the pixels in the input and output images will be the same).

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7 Data analysis

Astronomical datasets (images or tables) contain very valuable information, the tools in this section can help in analyzing, extracting, and quantifying that information. For example getting general or specific statistics of the dataset (with Statistics), detecting signal within a noisy dataset (with NoiseChisel), or creating a catalog from an input dataset (with MakeCatalog).

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7.1 Statistics

The distribution of values in a dataset can provide valuable information about it. For example, in an image, if it is a positively skewed distribution, we can see that there is significant data in the image. If the distribution is roughly symmetric, we can tell that there is no significant data in the image. In a table, when we need to select a sample of objects, it is important to first get a general view of the whole sample.

On the other hand, you might need to know certain statistical parameters of the dataset. For example, if we have run a detection algorithm on an image, and we want to see how accurate it was, one method is to calculate the average of the undetected pixels and see how reasonable it is (if detection is done correctly, the average of undetected pixels should be approximately equal to the background value, see Sky value). In a table, you might have calculated the magnitudes of a certain class of objects and want to get some general characteristics of the distribution immediately on the command-line (very fast!), to possibly change some parameters. The Statistics program is designed for such situations.

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7.1.1 Histogram and Cumulative Frequency Plot

Histograms and the cumulative frequency plots are both used to visually study the distribution of a dataset. A histogram shows the number of data points which lie within pre-defined intervals (bins). So on the horizontal axis we have the bin centers and on the vertical, the number of points that are in that bin. You can use it to get a general view of the distribution: which values have been repeated the most? how close/far are the most significant bins? Are there more values in the larger part of the range of the dataset, or in the lower part? Similarly, many very important properties about the dataset can be deduced from a visual inspection of the histogram. In the Statistics program, the histogram can be either output to a table to plot with your favorite plotting program119, or it can be shown with ASCII characters on the command-line, which is very crude, but good enough for a fast and on-the-go analysis, see the example in Invoking Statistics.

The width of the bins is only necessary parameter for a histogram. In the limiting case that the bin-widths tend to zero (while assuming the number of points in the dataset tend to infinity), then the histogram will tend to the probability density function of the distribution. When the absolute number of points in each bin is not relevant to the study (only the shape of the histogram is important), you can normalize a histogram so like the probability density function, the sum of all its bins will be one.

In the cumulative frequency plot of a distribution, the horizontal axis is the sorted data values and the y axis is the index of each data in the sorted distribution. Unlike a histogram, a cumulative frequency plot does not involve intervals or bins. This makes it less prone to any sort of bias or error that a given bin-width would have on the analysis. When a larger number of the data points have roughly the same value, then the cumulative frequency plot will become steep in that vicinity. This occurs because on the horizontal axis, there is little change while on the vertical axis, the indexes constantly increase. Normalizing a cumulative frequency plot means to divide each index (y axis) by the total number of data points (or the last value).

Unlike the histogram which has a limited number of bins, ideally the cumulative frequency plot should have one point for every data element. Even in small datasets (for example a \(200\times200\) image) this will result in an unreasonably large number of points to plot (40000)! As a result, for practical reasons, it is common to only store its value on a certain number of points (intervals) in the input range rather than the whole dataset, so you should determine the number of bins you want when asking for a cumulative frequency plot. In Gnuastro (and thus the Statistics program), the number reported for each bin is the total number of data points until the larger interval value for that bin. You can see an example histogram and cumulative frequency plot of a single dataset under the --asciihist and --asciicfp options of Invoking Statistics.

So as a summary, both the histogram and cumulative frequency plot in Statistics will work with bins. Within each bin/interval, the lower value is considered to be within then bin (it is inclusive), but its larger value is not (it is exclusive). Formally, an interval/bin between a and b is represented by [a, b). When the over-all range of the dataset is specified (with the --greaterequal, --lessthan, or --qrange options), the acceptable values of the dataset are also defined with a similar inclusive-exclusive manner. But when the range is determined from the actual dataset (none of these options is called), the last element in the dataset is included in the last bin’s count.

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7.1.2 Sigma clipping

Let’s assume that you have pure noise (centered on zero) with a clear Gaussian distribution, or see Photon counting noise. Now let’s assume you add very bright objects (signal) on the image which have a very sharp boundary. By a sharp boundary, we mean that there is a clear cutoff (from the noise) at the pixels the objects finish. In other words, at their boundaries, the objects do not fade away into the noise. In such a case, when you plot the histogram (see Histogram and Cumulative Frequency Plot) of the distribution, the pixels relating to those objects will be clearly separate from pixels that belong to parts of the image that did not have any signal (were just noise). In the cumulative frequency plot, after a steady rise (due to the noise), you would observe a long flat region were for a certain range of data (horizontal axis), there is no increase in the index (vertical axis).

Outliers like the example above can significantly bias the measurement of noise statistics. \(\sigma\)-clipping is defined as a way to avoid the effect of such outliers. In astronomical applications, cosmic rays (when they collide at a near normal incidence angle) are a very good example of such outliers. The tracks they leave behind in the image are perfectly immune to the blurring caused by the atmosphere and the aperture. They are also very energetic and so their borders are usually clearly separated from the surrounding noise. So \(\sigma\)-clipping is very useful in removing their effect on the data. See Figure 15 in Akhlaghi and Ichikawa, 2015.

\(\sigma\)-clipping is defined as the very simple iteration below. In each iteration, the range of input data might decrease and so when the outliers have the conditions above, the outliers will be removed through this iteration. The exit criteria will be discussed below.

  1. Calculate the standard deviation (\(\sigma\)) and median (\(m\)) of a distribution.
  2. Remove all points that are smaller or larger than \(m\pm\alpha\sigma\).
  3. Go back to step 1, unless the selected exit criteria is reached.

The reason the median is used as a reference and not the mean is that the mean is too significantly affected by the presence of outliers, while the median is less affected, see Quantifying signal in a tile. As you can tell from this algorithm, besides the condition above (that the signal have clear high signal to noise boundaries) \(\sigma\)-clipping is only useful when the signal does not cover more than half of the full data set. If they do, then the median will lie over the outliers and \(\sigma\)-clipping might remove the pixels with no signal.

There are commonly two exit criteria to stop the \(\sigma\)-clipping iteration:

When working on astronomical images, objects like galaxies and stars are blurred by the atmosphere and the telescope aperture, therefore their signal sinks into the noise very gradually. Galaxies in particular do not appear to have a clear high signal to noise cutoff at all. Therefore \(\sigma\)-clipping will not be useful in removing their effect on the data.

To gauge if \(\sigma\)-clipping will be useful for your dataset, look at the histogram (see Histogram and Cumulative Frequency Plot). The ASCII histogram that is printed on the command-line with --asciihist is good enough in most cases.

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7.1.3 Sky value

One of the most important aspects of a dataset is its reference value: the value of the dataset where there is no signal. Without knowing, and thus removing the effect of, this value it is impossible to compare the derived results of many high-level analyses over the dataset with other datasets (in the attempt to associate our results with the “real” world). In astronomy, this reference value is known as the “Sky” value: the value where there is no signal from objects (for example galaxies, stars, planets or comets). Depending on the dataset, the Sky value maybe a fixed value over the whole dataset, or it may vary based on location. For an example of the latter case, see Figure 11 in Akhlaghi and Ichikawa (2015).

Because of the significance of the Sky value in astronomical data analysis, we have devoted this subsection to it for a thorough review. We start with a thorough discussion on its definition (Sky value definition). In the astronomical literature, researchers use a variety of methods to estimate the Sky value, so in Sky value misconceptions) we review those and discuss their biases. From the definition of the Sky value, the most accurate way to estimate the Sky value is to run a detection algorithm (for example NoiseChisel) over the dataset and use the undetected pixels. However, there is also a more crude method that maybe useful when good direct detection is not initially possible (for example due to too many cosmic rays in a shallow image). A more crude (but simpler method) that is usable in such situations is discussed in Quantifying signal in a tile.

Next: , Previous: , Up: Sky value   [Contents][Index] Sky value definition

This analysis is taken from Akhlaghi and Ichikawa (2015). Let’s assume that all instrument defects – bias, dark and flat – have been corrected and the brightness (see Flux Brightness and magnitude) of a detected object, \(O\), is desired. The sources of flux on pixel \(i\)120 of the image can be written as follows:

The total flux in this pixel (\(T_i\)) can thus be written as:


By definition, \(D_i\) is detected and it can be assumed that it is correctly estimated (deblended) and subtracted, thus \(D_i=0\). There are also methods to detect and remove cosmic rays, for example the method described in van Dokkum (2001)121, or by comparing multiple exposures. This allows us to set \(C_i=0\). Note that in practice, \(D_i\) and \(U_i\) are correlated, because they both directly depend on the detection algorithm and its input parameters. Also note that no detection or cosmic ray removal algorithm is perfect. With these limitations in mind, the observed Sky value for this pixel (\(S_i\)) can be defined as


Therefore, as the detection process (algorithm and input parameters) becomes more accurate, or \(U_i\to0\), the sky value will tend to the background value or \(S_i\to B_i\). Therefore, while \(B_i\) is an inherent property of the data (pixel in an image), \(S_i\) depends on the detection process. Over a group of pixels, for example in an image or part of an image, this equation translates to the average of undetected pixels. With this definition of sky, the object flux in the data can be calculated with

$$T_{i}=S_{i}+O_{i} \quad\rightarrow\quad O_{i}=T_{i}-S_{i}.$$

Hence, the more accurately \(S_i\) is measured, the more accurately the brightness (sum of pixel values) of the target object can be measured (photometry). Any under-(over-)estimation in the sky will directly translate to an over-(under-)estimation of the measured object’s brightness. In the fainter outskirts of an object a very small fraction of the photo-electrons in the pixels actually belong to objects (see Figure 1b in Akhlaghi and Ichikawa (2015)). Therefore even a small over estimation of the sky value will result in the loss of a very large portion of most galaxies. Besides the lost area/brightness, this will also cause an over-estimation of the Sky value and thus even more under-estimation of the object’s brightness. It is thus very important to detect the diffuse flux of a target, even if they are not your primary target.

The Sky value is only correctly found when all the detected objects (\(D_i\) and \(C_i\)) have been removed from the data.

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As defined in Sky value, the sky value is only accurately defined when the detection algorithm is not significantly reliant on the sky value. In particular its detection threshold. However, most signal-based detection tools122 use the sky value as a reference to define the detection threshold. So these old techniques had to rely on approximations based on other assumptions about the data. A review of those other techniques can be seen in Appendix A of Akhlaghi and Ichikawa (2015)123. Since they were extensively used in astronomical data analysis for several decades, such approximations have given rise to a lot of misconceptions, ambiguities and disagreements about the sky value and how to measure it. As a summary, the major methods used until now were an approximation of the mode of the image pixel distribution and \(\sigma\)-clipping.

As discussed in Sky value, the sky value can only be correctly defined as the average of undetected pixels. Therefore all such approaches that try to approximate the sky value prior to detection are ultimately poor approximations.

Previous: , Up: Sky value   [Contents][Index] Quantifying signal in a tile

Put simply, noise can characterized with a certain spread about a characteristic value. In the Gaussian distribution (most commonly used to model noise) the spread is defined by the standard deviation about the characteristic mean. Before continuing let’s clarify some definitions first: Data is defined as the combination of signal and noise (so a noisy image is one data-set). Signal is defined as the mean of the noise on each element (after sky subtraction, see Sky value definition).

Let’s assume that the background (see Sky value definition) is subtracted and is zero. When a data set doesn’t have any signal (only noise), the mean, median and mode of the distribution are equal within statistical errors and approximately equal to the background value. Signal always has a positive value and will never become negative, see Figure 1 in Akhlaghi and Ichikawa (2015). Therefore, as more signal is added to the raw noise, the mean, median and mode of the dataset (which has both signal and noise) shift to the positive. The mean’s shift is the largest. The median shifts less, since it is defined based on an ordered distribution and so is not affected by a small number of outliers. The distribution’s mode shifts the least to the positive.

Inverting the argument above gives us a robust method to quantify the significance of signal in a dataset. Namely, when the mode and median of a distribution are approximately equal, we can argue that there is no significant signal. To allow for gradients (which are commonly present in ground-based images), we can consider the image to be made of a grid of tiles (see Tessellation124). Hence, from the difference of the mode and median on each tile, we can ‘detect’ the significance of signal in it. The median of a distribution is defined to be the value of the distribution’s middle point after sorting (or 0.5 quantile). Thus, to estimate the presence of signal, we’ll compare with the quantile of the mode with 0.5, if the difference is larger than the value given to the --modmedqdiff option, this tile will be ignored. You can read this option as “mode-median-quantile-diff”.

This method to use the input’s skewness is possible because of a new algorithm to find the mode of a distribution that was defined in Appendix C of Akhlaghi and Ichikawa (2015). However, the raw dataset’s distribution is noisy (noise also affects the sorting), so using the argument above on the raw input will give a noisy result. To decrease the noise/error in estimating the mode, we will use convolution (see Convolution process). Convolution decreases the range of the dataset and enhances its skewness, See Section 3.1.1 and Figure 4 in Akhlaghi and Ichikawa (2015). This enhanced skewness can be interpreted as an increase in the Signal to noise ratio of the objects buried in the noise. Therefore, to obtain an even better measure of the presence of signal in a mesh, the image can be convolved with a given kernel first.

Note that through the difference of the mode and median we have actually ‘detected’ data in the distribution. However this “detection” was only based on the total distribution of the data in each tile (a much lower resolution). This is the main limitation of this technique. The best approach is thus to do detection over the dataset, mask all the detected pixels and use the undetected regions to estimate the sky and its standard deviation.

The mean value of the tiles that have an approximately equal mode and median will be the Sky value. However there is one final hurdle: astronomical datasets are commonly plagued with Cosmic rays. Images of Cosmic rays aren’t smoothed by the atmosphere or telescope aperture, so they have sharp boundaries. Also, since they don’t occupy too many pixels, they don’t affect the mode and median calculation. But their very high values can greatly bias the calculation of the mean (recall how the mean shifts the fastest in the presence of outliers), see Figure 15 in Akhlaghi and Ichikawa (2015) for one example.

The effect of outliers like cosmic rays on the mean and standard deviation can be removed through \(\sigma\)-clipping, see Sigma clipping for a complete explanation. Therefore, after asserting that the mode and median are approximately equal in a tile (see Tessellation), the final Sky value and its standard deviation are determined after \(\sigma\)-clipping with the --sigmaclip option.

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7.1.4 Invoking Statistics

Statistics will print statistical measures of an input dataset (table column or image). The executable name is aststatistics with the following general template

$ aststatistics [OPTION ...] InputImage.fits

One line examples:

## Print some general statistics of input image:
$ aststatistics image.fits

## Print some general statistics of column named MAG_F160W:
$ aststatistics catalog.fits -h1 --column=MAG_F160W

## Make the histogram of the column named MAG_F160W:
$ aststatistics table.fits -cMAG_F160W --histogram

## Find the Sky value on image with a given kernel:
$ aststatistics image.fits --sky --kernel=kernel.fits

## Print Sigma-clipped results of records with a MAG_F160W
## column value between 26 and 27:
$ aststatistics cat.fits -cMAG_F160W -g26 -l27 --sigmaclip=3,0.2

## Print the median value of all records in column MAG_F160W that
## have a value larger than 3 in column PHOTO_Z:
$ aststatistics tab.txt -rPHOTO_Z -g3 -cMAG_F160W --median

An input image or table is necessary when processing is to be done. If any output file is to be created, the value to the --output option, is used as the base name for the generated files. Without --output, the input name will be used to generate an output name, see Automatic output. The options described below are particular to Statistics, but for general operations, it shares a large collection of options with the other Gnuastro programs, see Common options for the full list. Options can also be given in configuration files, for more, please see Configuration files.

The input dataset may have blank values (see Blank pixels), in this case, all blank pixels are ignored during the calculation. Initially, the full dataset will be read, but it is possible to select a specific range of data elements to use in the analysis of each run. You can either directly specify a minimum and maximum value for the range of data elements to use (with --greaterequal or --lessthan), or specify the range using quantiles (with --qrange). If a range is specified, all pixels outside of it are ignored before any processing.

The following set of options are for specifying the input/outputs of Statistics. There are many other input/output options that are common to all Gnuastro programs including Statistics, see Input/Output options for those.


The input column selector when the input file is a table. See Selecting table columns for a full description of how to use this option. For more on how tables are read in Gnuastro, please see Tables.


The reference column selector when the input file is a table. When a reference column is given, the range options below will be applied to this column and only elements in the input column that have a reference value in the correct range will be used. In practice this option allows you to select a subset of the input column based on values in another (the reference) column. All the statistical calculations will be done on the selected input column, not the reference column.

-g FLT

Limit the range of inputs into those with values greater and equal to what is given to this option. None of the values below this value will be used in any of the processing steps below.

-l FLT

Limit the range of inputs into those with values less-than what is given to this option. None of the values greater or equal to this value will be used in any of the processing steps below.


Specify the range of usable inputs using the quantile. This option can take one or two quantiles to specify the range. When only one number is input (let’s call it \(Q\)), the range will be those values in the quantile range \(Q\) to \(1-Q\). So when only one value is given, it must be less than 0.5. When two values are given, the first is used as the lower quantile range and the second is used as the larger quantile range.

The quantile of a given element in a dataset is defined by the fraction of its index to the total number of values in the sorted input array. So the smallest and largest values in the dataset have a quantile of 0.0 and 1.0. The quantile is a very useful non-parametric (making no assumptions about the input) relative measure to specify a range. It can best be understood in terms of the cumulative frequency plot, see Histogram and Cumulative Frequency Plot. The quantile of each horizontal axis value in the cumulative frequency plot is the vertical axis value associate with it.

When no operation is requested, Statistics will print some general basic properties of the input dataset on the command-line like the example below (ran on one of the output images of make check125). This default behavior is designed to help give you a general feeling of how the data are distributed and help in narrowing down your analysis.

$ aststatistics convolve_spatial_scaled_noised.fits     \
                --greaterequal=9500 --lessthan=11000
Statistics (GNU Astronomy Utilities) X.X
Input: convolve_spatial_scaled_noised.fits (hdu: 0)
Range: from (inclusive) 9500, upto (exclusive) 11000.
Unit: Brightness
  Number of elements:                      9074
  Minimum:                                 9622.35
  Maximum:                                 10999.7
  Mode:                                    10055.45996
  Mode quantile:                           0.4001983908
  Median:                                  10093.7
  Mean:                                    10143.98257
  Standard deviation:                      221.80834
 |                   **
 |                 ******
 |                 *******
 |                *********
 |              *************
 |              **************
 |            ******************
 |            ********************
 |          *************************** *
 |        ***************************************** ***
 |*  **************************************************************

Gnuastro’s Statistics is a very general purpose program, so to be able to easily understand this diversity in its operations (and how to possibly run them together), we’ll divided the operations into two types: those that don’t respect the position of the elements and those that do (by tessellating the input on a tile grid, see Tessellation). The former treat the whole dataset as one and can re-arrange all the elements (for example sort them), but the former do their processing on each tile independently. First, we’ll review the operations that work on the whole dataset.

The group of options below can be used to get single value measurement(s) of the whole dataset. They will print only the requested value as one field in a line/row, like the --mean, --median options. These options can be called any number of times and in any order. The outputs of all such options will be printed on one line following each other (with a space character between them). This feature makes these options very useful in scripts, or to redirect into programs like GNU AWK for higher-level processing. These are some of the most basic measures, Gnuastro is still under heavy development and this list will grow. If you want another statistical parameter, please contact us and we will do out best to add it to this list, see Suggest new feature.


Print the number of all used (non-blank and in range) elements.


Print the minimum value of all used elements.


Print the maximum value of all used elements.


Print the sum of all used elements.


Print the mean (average) of all used elements.


Print the standard deviation of all used elements.


Print the median of all used elements.

-u FLT[,FLT[,...]]

Print the values at the given quantiles of the input dataset. Any number of quantiles may be given and one number will be printed for each. Values can either be written as a single number or as fractions, but must be between zero and one (inclusive). Hence, in effect --quantile=0.25 --quantile=0.75 is equivalent to --quantile=0.25,3/4, or -u1/4,3/4.

The returned value is one of the elements from the dataset. Taking \(q\) to be your desired quantile, and \(N\) to be the total number of used (non-blank and within the given range) elements, the returned value is at the following position in the sorted array: \(round(q\times{}N\)).


Print the quantiles of the given values in the dataset. This option is the inverse of the --quantile and operates similarly except that the acceptable values are within the range of the dataset, not between 0 and 1. Formally it is known as the “Quantile function”.

Since the dataset is not continuous this function will find the nearest element of the dataset and use its position to estimate the quantile function.


Print the mode of all used elements. The mode is found through the mirror distribution which is fully described in Appendix C of Akhlaghi and Ichikawa 2015. See that section for a full description.

This mode calculation algorithm is non-parametric, so when the dataset is not large enough (larger than about 1000 elements usually), or doesn’t have a clear mode it can fail. In such cases, this option will return a value of nan (for the floating point NaN value).

As described in that paper, the easiest way to assess the quality of this mode calculation method is to use it’s symmetricity (see --modesym below). A better way would be to use the --mirror option to generate the histogram and cumulative frequency tables for any given mirror value (the mode in this case) as a table. If you generate plots like those shown in Figure 21 of that paper, then your mode is accurate.


Print the quantile of the mode. You can get the actual mode value from the --mode described above. In many cases, the absolute value of the mode is irrelevant, but its position within the distribution is important. In such cases, this option will become handy.


Print the symmetricity of the calculated mode. See the description of --mode for more. This mode algorithm finds the mode based on how symmetric it is, so if the symmetricity returned by this option is too low, the mode is not too accurate. See Appendix C of Akhlaghi and Ichikawa 2015 for a full description. In practice, symmetricity values larger than 0.2 are mostly good.


Print the value in the distribution where the mirror and input distributions are no longer symmetric, see --mode and Appendix C of Akhlaghi and Ichikawa 2015 for more.

The list of options below are for those statistical operations that output more than one value. So while they can be called together in one run, their outputs will be distinct (each one’s output will usually be printed in more than one line).


Print an ASCII histogram of the usable values within the input dataset along with some basic information like the example below (from the UVUDF catalog126). The width and height of the histogram (in units of character widths and heights on your command-line terminal) can be set with the --numasciibins (for the width) and --asciiheight options.

For a full description of the histogram, please see Histogram and Cumulative Frequency Plot. An ASCII plot is certainly very crude and cannot be used in any publication, but it is very useful for getting a general feeling of the input dataset very fast and easily on the command-line without having to take your hands off the keyboard (which is a major distraction!). If you want to try it out, you can write it all in one line and ignore the \ and extra spaces.

$ aststatistics uvudf_rafelski_2015.fits.gz --hdu=1         \
                --column=MAG_F160W --lessthan=40            \
                --asciihist --numasciibins=55
ASCII Histogram:
Number: 8593
Y: (linear: 0 to 660)
X: (linear: 17.7735 -- 31.4679, in 55 bins)
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Print the cumulative frequency plot of the usable elements in the input dataset. Please see descriptions under --asciihist for more, the example below is from the same input table as that example. To better understand the cumulative frequency plot, please see Histogram and Cumulative Frequency Plot.

$ aststatistics uvudf_rafelski_2015.fits.gz --hdu=1         \
                --column=MAG_F160W --lessthan=40            \
                --asciicfp --numasciibins=55
ASCII Cumulative frequency plot:
Y: (linear: 0 to 8593)
X: (linear: 17.7735 -- 31.4679, in 55 bins)
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Save the histogram of the usable values in the input dataset into a table. The first column is the value at the center of the bin and the second is the number of points in that bin. If the --cumulative option is also called with this option in a run, then the table will have three columns (the third is the cumulative frequency plot). Through the --numbins, --onebinstart, or --manualbinrange, you can modify the first column values and with --normalize and --maxbinone you can modify the second columns. See below for the description of each.

By default (when no --output is specified) a plain text table will be created, see Gnuastro text table format. If a FITS name is specified, you can use the common option --tableformat to have it as a FITS ASCII or FITS binary format, see Common options. This table can then be fed into your favorite plotting tool and get a much more clean and nice histogram than what the raw command-line can offer you (with the --asciihist option).


Save the cumulative frequency plot of the usable values in the input dataset into a table, similar to --histogram.


Do \(\sigma\)-clipping on the usable pixels of the input dataset. See Sigma clipping for a full description on \(\sigma\)-clipping and also to better understand this option. The \(\sigma\)-clipping parameters can be set through the --sclipparams option (see below).


Make a histogram and cumulative frequency plot of the mirror distribution for the given dataset when the mirror is located at the value to this option. The mirror distribution is fully described in Appendix C of Akhlaghi and Ichikawa 2015 and currently it is only used to calculate the mode (see --mode).

Just note that the mirror distribution is a discrete distribution like the input, so while you may give any number as the value to this option, the actual mirror value is the closest number in the input dataset to this value. If the two numbers are different, Statistics will warn you of the actual mirror value used.

This option will make a table as output. Depending on your selected name for the output, it will be either a FITS table or a plain text table (which is the default). It contains three columns: the first is the center of the bins, the second is the histogram (with the largest value set to 1) and the third is the normalized cumulative frequency plot of the mirror distribution. The bins will be positioned such that the mode is on the starting interval of one of the bins to make it symmetric around the mirror. With this output file and the input histogram (that you can generate in another run of Statistics, using the --onebinvalue), it is possible to make plots like Figure 21 of Akhlaghi and Ichikawa 2015.

The list of options below allow customization of the histogram and cumulative frequency plots (for the --histogram, --cumulative, --asciihist, and --asciicfp options).


The number of bins (rows) to use in the histogram and the cumulative frequency plot tables (outputs of --histogram and --cumulative).


The number of bins (characters) to use in the ASCII plots when printing the histogram and the cumulative frequency plot (outputs of --asciihist and --asciicfp).


The number of lines to use when printing the ASCII histogram and cumulative frequency plot on the command-line (outputs of --asciihist and --asciicfp).


Normalize the histogram or cumulative frequency plot tables (outputs of --histogram and --cumulative). For a histogram, the sum of all bins will become one and for a cumulative frequency plot the last bin value will be one.


Divide all the histogram values by the maximum bin value so it becomes one and the rest are similarly scaled. In some situations (for example if you want to plot the histogram and cumulative frequency plot in one plot) this can be very useful.


Make sure that one bin starts with the value to this option. In practice, this will shift the bins used to find the histogram and cumulative frequency plot such that one bin’s lower interval becomes this value.

For example when a histogram range includes negative and positive values and zero has a special significance in your analysis, then zero might fall somewhere in one bin. As a result that bin will have counts of positive and negative. By setting --onebinstart=0, you can make sure that one bin will only count negative values in the vicinity of zero and the next bin will only count positive ones in that vicinity.

Note that by default, the first row of the histogram and cumulative frequency plot show the central values of each bin. So in the example above you will not see the 0.000 in the first column, you will see two symmetric values.

If the value is not within the usable input range, this option will be ignored. When it is, this option is the last operation before the bins are finalized, therefore it has a higher priority than options like --manualbinrange.


Use the values given to the --greaterequal and --lessthan to define the range of all bin-based calculations like the histogram. This option itself doesn’t take any value, but just tells the program to use the values of those two options instead of the minimum and maximum values of a plot. If any of the two options are not given, then the minimum or maximum will be used respectively. Therefore, if none of them are called calling this option is redundant.

The --onebinstart option has a higher priority than this option. In other words, --onebinstart takes effect after the range has been finalized and the initial bins have been defined, therefore it has the power to (possibly) shift the bins. If you want to manually set the range of the bins and have one bin on a special value, it is thus better to avoid --onebinstart.

All the options described until now were from the first class of operations discussed above: those that treat the whole dataset as one. However. It often happens that the relative position of the dataset elements over the dataset is significant. For example you don’t want one median value for the whole input image, you want to know how the median changes over the image. For such operations, the input has to be tessellated (see Tessellation). Thus this class of options can’t currently be called along with the options above in one run of Statistics.


Do the respective single-valued calculation over one tile of the input dataset, not the whole dataset. This option must be called with at least one of the single valued options discussed above (for example --mean or --quantile). The output will be a file in the same format as the input. If the --oneelempertile option is called, then one element/pixel will be used for each tile (see Processing options). Otherwise, the output will have the same size as the input, but each element will have the value corresponding to that tile’s value. If multiple single valued operations are called, then for each operation there will be one extension in the output FITS file.


Estimate the Sky value on each tile as fully described in Quantifying signal in a tile. As described in that section, several options are necessary to configure the Sky estimation which are listed below. The output file will have two extensions: the first is the Sky value and the second is the Sky standard deviation on each tile. Similar to --ontile, if the --oneelempertile option is called, then one element/pixel will be used for each tile (see Processing options).

The parameters for estimating the sky value can be set with the following options, except for the --sclipparams option (which is also used by the --sigmaclip), the rest are only used for the Sky value estimation.


File name of kernel to help in estimating the significance of signal in a tile, see Quantifying signal in a tile.


Kernel HDU to help in estimating the significance of signal in a tile, see Quantifying signal in a tile.


Maximum distance (as a multiple of error) to estimate the difference between the input and mirror distributions in finding the mode, see Appendix C of Akhlaghi and Ichikawa 2015, also see Quantifying signal in a tile.


The maximum acceptable distance between the mode and median, see Quantifying signal in a tile.


The \(\sigma\)-clipping parameters, see Sigma clipping. This option takes two values which are separated by a comma (,). Each value can either be written as a single number or as a fraction of two numbers (for example 3,1/10). The first value to this option is the multiple of \(\sigma\) that will be clipped (\(\alpha\) in that section). The second value is the exit criteria. If it is less than 1, then it is interpreted as tolerance and if it is larger than one it is a specific number. Hence, in the latter case the value must be an integer.


Width of a flat kernel to convolve the interpolated tile values. Tile interpolation is done using the median of the --interpnumngb neighbors of each tile (see Processing options). If this option is given a value of zero or one, no smoothing will be done. Without smoothing, strong boundaries will probably be created between the values estimated for each tile. It is thus good to smooth the interpolated image so strong discontinuities do not show up in the final Sky values. The smoothing is done through convolution (see Convolution process) with a flat kernel, so the value to this option must be an odd number.


Don’t set the input’s blank pixels to blank in the output Sky and Sky standard deviation datasets. This is only applicable when the output has the same size as the input, in other words, when --oneelempertile isn’t called. By default, blank values in the input (commonly on the edges which are outside the survey/field area) will be set to blank in the output Sky and Sky standard deviation also.


Create a multi-extension FITS file showing the steps that were used to estimate the Sky value over the input, see Quantifying signal in a tile. The file will have two extensions for each step (one for the Sky and one for the Sky standard deviation).

Next: , Previous: , Up: Data analysis   [Contents][Index]

7.2 NoiseChisel

Once instrumental signatures are removed from the raw data (image) in the initial reduction process (see Data manipulation). You are naturally eager to start answering the scientific questions that motivated the data collection in the first place. However, the raw dataset/image is just an array of values/pixels, that is all! These raw values cannot directly be used to answer your scientific questions: for example “how many galaxies are there in the image?”.

The first high-level step in your analysis of your targets will thus be to classify, or label, the dataset elements (pixels) into two classes: 1) noise, where random effects are the major contributor to the value, and 2) signal, where non-random factors (for example light from a distant galaxy) are present. This classification of the elements in a dataset is formally known as detection.

In an observational/experimental dataset, signal is always buried in noise: only mock/simulated datasets are free of noise. Therefore detection, or the process of separating signal from noise, determines the number of objects you study and the accuracy of any higher-level measurement you do on them. Detection is thus the most important step of any analysis and is not trivial. In particular, the most scientifically interesting astronomical targets are faint, can have a large variety of morphologies, along with a large distribution in brightness and size. Therefore when noise is significant, proper detection of your targets is a uniquely decisive step in your final scientific analysis/result.

NoiseChisel is Gnuastro’s program for detection of targets that don’t have a sharp border (almost all astronomical objects). When the targets have a sharp edges/border (for example cells in biological imaging), a simple threshold is enough to separate them from noise and each other (if they are not touching). To detect such sharp-edged targets, you can use Gnuastro’s Arithmetic program in a command like below (assuming the threshold is 100, see Arithmetic):

$ astarithmetic in.fits 100 gt 2 connected-components

Since almost no astronomical target has such sharp edges, we need a more advanced detection methodology. NoiseChisel uses a new noise-based paradigm for detection of very extended and diffuse targets that are drowned deeply in the ocean of noise. It was initially introduced in Akhlaghi and Ichikawa [2015]. The name of NoiseChisel is derived from the first thing it does after thresholding the dataset: to erode it. In mathematical morphology, erosion on pixels can be pictured as carving-off boundary pixels. Hence, what NoiseChisel does is similar to what a wood chisel or stone chisel do. It is just not a hardware, but a software. In fact, looking at it as a chisel and your dataset as a solid cube of rock will greatly help in effectively understanding and optimally using it: with NoiseChisel you literally carve your targets out of the noise. Try running it with the --checkdetection option to see each step of the carving process on your input dataset.

NoiseChisel’s primary output is a binary detection map with the same size as the input but only with two values: 0 and 1. Pixels that don’t harbor any detected signal (noise) are given a label (or value) of zero and those with a value of 1 have been identified as hosting signal.

Segmentation is the process of classifying the signal into higher-level constructs. For example if you have two separate galaxies in one image, by default NoiseChisel will give a value of 1 to the pixels of both, but after segmentation, the pixels in each will get separate labels. NoiseChisel is only focused on detection (separating signal from noise), to segment the signal (into separate galaxies for example), Gnuastro has a separate specialized program Segment. NoiseChisel’s output can be directly/readily fed into Segment.

For more on NoiseChisel’s output format and its benefits (especially in conjunction with Segment and later MakeCatalog), please see Akhlaghi [2016]. Just note that when that paper was published, Segment was not yet spun-off into a separate program, and NoiseChisel done both detection and segmentation.

NoiseChisel’s output is designed to be generic enough to be easily used in any higher-level analysis. If your targets are not touching after running NoiseChisel and you aren’t interested in their sub-structure, you don’t need the Segment program at all. You can ask NoiseChisel to find the connected pixels in the output with the --label option. In this case, the output won’t be a binary image any more, the signal will have counters/labels starting from 1 for each connected group of pixels. You can then directly feed NoiseChisel’s output into MakeCatalog for measurements over the detections and the production of a catalog (see MakeCatalog).

Thanks to the published papers mentioned above, there is no need to provide a more complete introduction to NoiseChisel in this book. However, published papers cannot be updated any more, but the software has evolved/changed. The changes since publication are documented in NoiseChisel changes after publication. Afterwards, in Invoking NoiseChisel, the details of running NoiseChisel and its options are discussed.

As discussed above, detection is one of the most important steps for your scientific result. It is therefore very important to obtain a good understanding of NoiseChisel (and afterwards Segment and MakeCatalog). We thus strongly recommend that after reading the papers above and the respective sections of Gnuastro’s book, you play a little with the settings (in the order presented in the paper and Invoking NoiseChisel) on a dataset you are familiar with and inspect all the check images (options starting with --check) to see the effect of each parameter.

We strongly recommend going over the two tutorials of General program usage tutorial and Detecting large extended targets. They are designed to show how to most effectively use NoiseChisel for the detection of small faint objects and large extended objects. In the meantime, they will show you the modular principle behind Gnuastro’s programs and how they are built to complement, and build upon, each other. General program usage tutorial culminates in using NoiseChisel to detect galaxies and use its outputs to find the galaxy colors. Defining colors is a very common process in most science-cases. Therefore it is also recommended to (patiently) complete that tutorial for optimal usage of NoiseChisel in conjunction with all the other Gnuastro programs. Detecting large extended targets shows you can optimize NoiseChisel’s settings for very extended objects to successfully carve out to signal-to-noise ratio levels of below 1/10.

In NoiseChisel changes after publication, we’ll review the changes in NoiseChisel since the publication of Akhlaghi and Ichikawa [2015]. We will then review NoiseChisel’s input, detection, and output options in NoiseChisel input, Detection options, and NoiseChisel output.

Next: , Previous: , Up: NoiseChisel   [Contents][Index]

7.2.1 NoiseChisel changes after publication

NoiseChisel was initially introduced in Akhlaghi and Ichikawa [2015]. It is thus strongly recommended to read this paper for a good understanding of what it does and how each parameter influences the output. To help in understanding how it works, that paper has a large number of figures showing every step on multiple mock and real examples.

However, the paper cannot be updated anymore, but NoiseChisel has evolved (and will continue to do so): better algorithms or steps have been found, thus options will be added or removed. This book is thus the final and definitive guide to NoiseChisel. The aim of this section is to make the transition from the paper to the installed version on your system, as smooth as possible with the list below. For a more detailed list of changes in previous Gnuastro releases/versions, please see the NEWS file127.

The most important change since the publication of that paper is that from Gnuastro 0.6, NoiseChisel is only in charge on detection. Segmentation of the detected signal was spun-off into a separate program: Segment. This spin-off allows much greater creativity and is in the spirit of Gnuastro’s modular design (see Program design philosophy). Below you can see the major changes since that paper was published. First, the removed options/features are discussed, then we review the new features that have been added.

Removed features/options:

Added options:

Previous: , Up: NoiseChisel   [Contents][Index]

7.2.2 Invoking NoiseChisel

NoiseChisel will detect signal in noise producing a multi-extension dataset containing a binary detection map which is the same size as the input. Its output can be readily used for input into Segment, for higher-level segmentation, or MakeCatalog to do measurements and generate a catalog. The executable name is astnoisechisel with the following general template

$ astnoisechisel [OPTION ...] InputImage.fits

One line examples:

## Detect signal in input.fits.
$ astnoisechisel input.fits

## Inspect all the detection steps after changing a parameter.
$ astnoisechisel input.fits --qthresh=0.4 --checkdetection

## Detect signal assuming input has 4 amplifier channels along first
## dimension and 1 along the second. Also set the regular tile size
## to 100 along both dimensions:
$ astnoisechisel --numchannels=4,1 --tilesize=100,100 input.fits

If NoiseChisel is to do processing (for example you don’t want to get help, or see the values to each input parameter), an input image should be provided with the recognized extensions (see Arguments). NoiseChisel shares a large set of common operations with other Gnuastro programs, mainly regarding input/output, general processing steps, and general operating modes. To help in a unified experience between all of Gnuastro’s programs, these operations have the same command-line options, see Common options for a full list/description (they are not repeated here).

As in all Gnuastro programs, options can also be given to NoiseChisel in configuration files. For a thorough description on Gnuastro’s configuration file parsing, please see Configuration files. All of NoiseChisel’s options with a short description are also always available on the command-line with the --help option, see Getting help. To inspect the option values without actually running NoiseChisel, append your command with --printparams (or -P).

NoiseChisel’s input image may contain blank elements (see Blank pixels). Blank elements will be ignored in all steps of NoiseChisel. Hence if your dataset has bad pixels which should be masked with a mask image, please use Gnuastro’s Arithmetic program (in particular its where operator) to convert those pixels to blank pixels before running NoiseChisel. Gnuastro’s Arithmetic program has bitwise operators helping you select specific kinds of bad-pixels when necessary.

A convolution kernel can also be optionally given. If a value (file name) is given to --kernel on the command-line or in a configuration file (see Configuration files), then that file will be used to convolve the image prior to thresholding. Otherwise a default kernel will be used. The default kernel is a 2D Gaussian with a FWHM of 2 pixels truncated at 5 times the FWHM. This choice of the default kernel is discussed in Section 3.1.1 of Akhlaghi and Ichikawa [2015]. See Convolution kernel for kernel related options. Passing none to --kernel will disable convolution. On the other hand, through the --convolved option, you may provide an already convolved image, see descriptions below for more.

NoiseChisel defines two tessellations over the input (see Tessellation). This enables it to deal with possible gradients in the input dataset and also significantly improve speed by processing each tile on different threads simultaneously. Tessellation related options are discussed in Processing options. In particular, NoiseChisel uses two tessellations (with everything between them identical except the tile sizes): a fine-grained one with smaller tiles (used in thresholding and Sky value estimations) and another with larger tiles which is used for pseudo-detections over non-detected regions of the image. The common Tessellation options described in Processing options define all parameters of both tessellations. The large tile size for the latter tessellation is set through the --largetilesize option. To inspect the tessellations on your input dataset, run NoiseChisel with --checktiles.

Usage TIP: Frequently use the options starting with --check. Since the noise properties differ between different datasets, you can often play with the parameters/options for a better result than the default parameters. You can start with --checkdetection for the main steps. For the full list of NoiseChisel’s checking options please run:

$ astnoisechisel --help | grep check

Below, we’ll discuss NoiseChisel’s options, classified into two general classes, to help in easy navigation. NoiseChisel input mainly discusses the basic options relating to inputs and prior to the detection process detection. Afterwards, Detection options fully describes every configuration parameter (option) related to detection and how they affect the final result. The order of options in this section follow the logical order within NoiseChisel. On first reading (while you are still new to NoiseChisel), it is therefore strongly recommended to read the options in the given order below. The output of --printparams (or -P) also has this order. However, the output of --help is sorted alphabetically. Finally, in NoiseChisel output the format of NoiseChisel’s output is discussed.

Next: , Previous: , Up: Invoking astnoisechisel   [Contents][Index] NoiseChisel input

The options here can be used to configure the inputs and output of NoiseChisel, along with some general processing options. Recall that you can always see the full list of Gnuastro’s options with the --help (see Getting help), or --printparams (or -P) to see their values (see Operating mode options).

-k STR

File name of kernel to smooth the image before applying the threshold, see Convolution kernel. If no convolution is needed, give this option a value of none.

The first step of NoiseChisel is to convolve/smooth the image and use the convolved image in multiple steps including the finding and applying of the quantile threshold (see --qthresh).

The --kernel option is not mandatory. If not called, a 2D Gaussian profile with a FWHM of 2 pixels truncated at 5 times the FWHM is used. This choice of the default kernel is discussed in Section 3.1.1 of Akhlaghi and Ichikawa [2015]. You can use MakeProfiles to build a kernel with any of its recognized profile types and parameters. For more details, please see MakeProfiles output dataset. For example, the command below will make a Moffat kernel (with \(\beta=2.8\)) with FWHM of 2 pixels truncated at 10 times the FWHM.

$ astmkprof --oversample=1 --kernel=moffat,2,2.8,10

Since convolution can be the slowest step of NoiseChisel, for large datasets, you can convolve the image once with Gnuastro’s Convolve (see Convolve), and use the --convolved option to feed it directly to NoiseChisel. This can help getting faster results when you are playing/testing the higher-level options.


HDU containing the kernel in the file given to the --kernel option.


Use this file as the convolved image and don’t do convolution (ignore --kernel). NoiseChisel will just check the size of the given dataset is the same as the input’s size. If a wrong image (with the same size) is given to this option, the results (errors, bugs, and etc) are unpredictable. So please use this option with care and in a highly controlled environment, for example in the scenario discussed below.

In almost all situations, as the input gets larger, the single most CPU (and time) consuming step in NoiseChisel (and other programs that need a convolved image) is convolution. Therefore minimizing the number of convolutions can save a significant amount of time in some scenarios. One such scenario is when you want to segment NoiseChisel’s detections using the same kernel (with Segment, which also supports this --convolved option). This scenario would require two convolutions of the same dataset: once by NoiseChisel and once by Segment. Using this option in both programs, only one convolution (prior to running NoiseChisel) is enough.

Another common scenario where this option can be convenient is when you are testing NoiseChisel (or Segment) for the best parameters. You have to run NoiseChisel multiple times and see the effect of each change. However, once you are happy with the kernel, re-convolving the input on every change of higher-level parameters will greatly hinder, or discourage, further testing. With this option, you can convolve the input image with your chosen kernel once before running NoiseChisel, then feed it to NoiseChisel on each test run and thus save valuable time for better/more tests.

To build your desired convolution kernel, you can use MakeProfiles. To convolve the image with a given kernel you can use Convolve. Spatial domain convolution is mandatory: in the frequency domain, blank pixels (if present) will cover the whole image and gradients will appear on the edges, see Spatial vs. Frequency domain.

Below you can see an example of the second scenario: you want to see how variation of the growth level (through the --detgrowquant option) will affect the final result. Recall that you can ignore all the extra spaces, new lines, and backslash’s (‘\’) if you are typing in the terminal. In a shell script, remove the $ signs at the start of the lines.

## Make the kernel to convolve with.
$ astmkprof --oversample=1 --kernel=gaussian,2,5

## Convolve the input with the given kernel.
$ astconvolve input.fits --kernel=kernel.fits                \
              --domain=spatial --output=convolved.fits

## Run NoiseChisel with seven growth quantile values.
$ for g in 60 65 70 75 80 85 90; do                          \
    astnoisechisel input.fits --convolved=convolved.fits     \
                   --detgrowquant=0.$g --output=$g.fits;     \

The HDU/extension containing the convolved image in the file given to --convolved.

-w STR

File name of a wider kernel to use in estimating the difference of the mode and median in a tile (this difference is used to identify the significance of signal in that tile, see Quantifying signal in a tile). As displayed in Figure 4 of Akhlaghi and Ichikawa [2015], a wider kernel will help in identifying the skewness caused by data in noise. The image that is convolved with this kernel is only used for this purpose. Once the mode is found to be sufficiently close to the median, the quantile threshold is found on the image convolved with the sharper kernel (--kernel), see --qthresh).

Since convolution will significantly slow down the processing, this feature is optional. When it isn’t given, the image that is convolved with --kernel will be used to identify good tiles and apply the quantile threshold. This option is mainly useful in conditions were you have a very large, extended, diffuse signal that is still present in the usable tiles when using --kernel. See Detecting large extended targets for a practical demonstration on how to inspect the tiles used in identifying the quantile threshold.


HDU containing the kernel file given to the --widekernel option.


The size of each tile for the tessellation with the larger tile sizes. Except for the tile size, all the other parameters for this tessellation are taken from the common options described in Processing options. The format is identical to that of the --tilesize option that is discussed in that section.

Next: , Previous: , Up: Invoking astnoisechisel   [Contents][Index] Detection options

Detection is the process of separating the pixels in the image into two groups: 1) Signal and 2) Noise. Through the parameters below, you can customize the detection process in NoiseChisel. Recall that you can always see the full list of Gnuastro’s options with the --help (see Getting help), or --printparams (or -P) to see their values (see Operating mode options).

-r FLT

Maximum distance (as a multiple of error) to estimate the difference between the input and mirror distributions in finding the mode, see Appendix C of Akhlaghi and Ichikawa 2015, also see Quantifying signal in a tile.


The maximum acceptable distance between the mode and median, see Quantifying signal in a tile. The quantile threshold will be found on tiles that satisfy this mode and median difference.

-t FLT

The quantile threshold to apply to the convolved image. The detection process begins with applying a quantile threshold to each of the tiles in the small tessellation. The quantile is only calculated for tiles that don’t have any significant signal within them, see Quantifying signal in a tile. Interpolation is then used to give a value to the unsuccessful tiles and it is finally smoothed.

The quantile value is a floating point value between 0 and 1. Assume that we have sorted the \(N\) data elements of a distribution (the pixels in each mesh on the convolved image). The quantile (\(q\)) of this distribution is the value of the element with an index of (the nearest integer to) \(q\times{N}\) in the sorted data set. After thresholding is complete, we will have a binary (two valued) image. The pixels above the threshold are known as foreground pixels (have a value of 1) while those which lie below the threshold are known as background (have a value of 0).


Only keep tiles which have a q-thresh value above the given quantile of the all successful tiles. Hence, when given a value of 1, this option will be ignored. When there is more than one channel (and --workoverch is not called), quantile calculation and application will be done on each channel independently.

This option is useful when a large and diffuse (almost flat within each tile) signal exists with very small regions of Sky. The flatness of the profile will cause it to successfully pass the tests of Quantifying signal in a tile. As a result, without this option the flat and diffuse signal will be interpreted as sky. In such cases, you can see the status of the tiles with the --checkqthresh option (first image extension is enough) and select a quantile through this option to ignore the measured values in the higher-valued tiles.

This option can also be useful when there are large bright objects in the image with large flat wings which can also pass the tests and give outlier values. When there is a sky gradient over the image (mainly due to post-processing issues like bad flat fielding), this option must be set to 1 so it is completely ignored and the sky gradient is accurately measured and subtracted.

To get an estimate of the measured qthresh distribution, you can run the following commands. The first will create the qthresh check image with one value/pixel per tile (see Processing options). Open the image in a FITS viewer and inspect it. The second command below will print the basic information about the distribution of values and the third will print the value at the 0.4 quantile. Recall that Gnuastro’s Statistics program ignores blank values (in this case: tiles with significant signal), see Statistics.

$ astnoisechisel image.fits --checkqthresh --oneelempertile
$ aststatistics image_qthresh.fits
$ aststatistics image_qthresh.fits --quantile=0.4

Width of flat kernel used to smooth the interpolated quantile thresholds, see --qthresh for more.


Check the quantile threshold values on the mesh grid. A file suffixed with _qthresh.fits will be created showing each step. With this option, NoiseChisel will abort as soon as quantile estimation has been completed, allowing you to inspect the steps leading to the final quantile threshold, this can be disabled with --continueaftercheck. By default the output will have the same pixel size as the input, but with the --oneelempertile option, only one pixel will be used for each tile (see Processing options).

-e INT

The number of erosions to apply to the binary thresholded image. Erosion is simply the process of flipping (from 1 to 0) any of the foreground pixels that neighbor a background pixel. In a 2D image, there are two kinds of neighbors, 4-connected and 8-connected neighbors. You can specify which type of neighbors should be used for erosion with the --erodengb option, see below.

Erosion has the effect of shrinking the foreground pixels. To put it another way, it expands the holes. This is a founding principle in NoiseChisel: it exploits the fact that with very low thresholds, the holes in the very low surface brightness regions of an image will be smaller than regions that have no signal. Therefore by expanding those holes, we are able to separate the regions harboring signal.


The type of neighborhood (structuring element) used in erosion, see --erode for an explanation on erosion. Only two integer values are acceptable: 4 or 8. In 4-connectivity, the neighbors of a pixel are defined as the four pixels on the top, bottom, right and left of a pixel that share an edge with it. The 8-connected neighbors on the other hand include the 4-connected neighbors along with the other 4 pixels that share a corner with this pixel. See Figure 6 (a) and (b) in Akhlaghi and Ichikawa (2015) for a demonstration.


Pure erosion is going to carve off sharp and small objects completely out of the detected regions. This option can be used to avoid missing such sharp and small objects (which have significant pixels, but not over a large area). All pixels with a value larger than the significance level specified by this option will not be eroded during the erosion step above. However, they will undergo the erosion and dilation of the opening step below.

Like the --qthresh option, the significance level is determined using the quantile (a value between 0 and 1). Just as a reminder, in the normal distribution, \(1\sigma\), \(1.5\sigma\), and \(2\sigma\) are approximately on the 0.84, 0.93, and 0.98 quantiles.

-p INT

Depth of opening to be applied to the eroded binary image. Opening is a composite operation. When opening a binary image with a depth of \(n\), \(n\) erosions (explained in --erode) are followed by \(n\) dilations. Simply put, dilation is the inverse of erosion. When dilating an image any background pixel is flipped (from 0 to 1) to become a foreground pixel. Dilation has the effect of fattening the foreground. Note that in NoiseChisel, the erosion which is part of opening is independent of the initial erosion that is done on the thresholded image (explained in --erode). The structuring element for the opening can be specified with the --openingngb option. Opening has the effect of removing the thin foreground connections (mostly noise) between separate foreground ‘islands’ (detections) thereby completely isolating them. Once opening is complete, we have initial detections.


The structuring element used for opening, see --erodengb for more information about a structuring element.


The \(\sigma\)-clipping parameters, see Sigma clipping. This option takes two values which are separated by a comma (,). Each value can either be written as a single number or as a fraction of two numbers (for example 3,1/10). The first value to this option is the multiple of \(\sigma\) that will be clipped (\(\alpha\) in that section). The second value is the exit criteria. If it is less than 1, then it is interpreted as tolerance and if it is larger than one it is assumed to be the fixed number of iterations. Hence, in the latter case the value must be an integer.


Minimum fraction (value between 0 and 1) of Sky (undetected) areas in a tile. Only tiles with a fraction of undetected pixels (Sky) larger than this value will be used to estimate the Sky value. NoiseChisel uses this option value twice to estimate the Sky value: after initial detections and in the end when false detections have been removed.

Because of the PSF and their intrinsic amorphous properties, astronomical objects (except cosmic rays) never have a clear cutoff and commonly sink into the noise very slowly. Even below the very low thresholds used by NoiseChisel. So when a large fraction of the area of one mesh is covered by detections, it is very plausible that their faint wings are present in the undetected regions (hence causing a bias in any measurement). To get an accurate measurement of the above parameters over the tessellation, tiles that harbor too many detected regions should be excluded. The used tiles are visible in the respective --check option of the given step.


Check the initial approximation of the sky value and its standard deviation in a FITS file ending with _detsky.fits. With this option, NoiseChisel will abort as soon as the sky value used for defining pseudo-detections is complete. This allows you to inspect the steps leading to the final quantile threshold, this behavior can be disabled with --continueaftercheck. By default the output will have the same pixel size as the input, but with the --oneelempertile option, only one pixel will be used for each tile (see Processing options).


The detection threshold: a multiple of the initial sky standard deviation added with the initial sky approximation (which you can inspect with --checkdetsky). This flux threshold is applied to the initially undetected regions on the unconvolved image. The background pixels that are completely engulfed in a 4-connected foreground region are converted to background (holes are filled) and one opening (depth of 1) is applied over both the initially detected and undetected regions. The Signal to noise ratio of the resulting ‘psudo-detections’ are used to identify true vs. false detections. See Section 3.1.5 and Figure 7 in Akhlaghi and Ichikawa (2015) for a very complete explanation.

-m INT

The minimum area to calculate the Signal to noise ratio on the psudo-detections of both the initially detected and undetected regions. When the area in a psudo-detection is too small, the Signal to noise ratio measurements will not be accurate and their distribution will be heavily skewed to the positive. So it is best to ignore any psudo-detection that is smaller than this area. Use --detsnhistnbins to check if this value is reasonable or not.


Save the S/N values of the pseudo-detections and dilated detections into three files ending with _detsn_sky.XXX, _detsn_det.XXX, and _detsn_dilated.XXX. The .XXX is determined from the --tableformat option (see Input/Output options, for example .txt or .fits). You can use these to inspect the S/N values and their distribution (in combination with the --checkdetection option to see where the pseudo-detections are). You can use Gnuastro’s Statistics to make a histogram of the distribution or any other analysis you would like for better understanding of the distribution (for example through a histogram).

With this option, NoiseChisel will abort as soon as the tables are created. This allows you to inspect the steps leading to the final quantile threshold, this behavior (to abort NoiseChisel) can be disabled with --continueaftercheck.


The minimum number of ‘pseudo-detections’ over the undetected regions to identify a Signal-to-Noise ratio threshold. The Signal to noise ratio (S/N) of false pseudo-detections in each tile is found using the quantile of the S/N distribution of the psudo-detections over the undetected pixels in each mesh. If the number of S/N measurements is not large enough, the quantile will not be accurate (can have large scatter). For example if you set --snquant=0.99 (or the top 1 percent), then it is best to have at least 100 S/N measurements.

-c FLT

The quantile of the Signal to noise ratio distribution of the psudo-detections in each mesh to use for filling the large mesh grid. Note that this is only calculated for the large mesh grids that satisfy the minimum fraction of undetected pixels (value of --minbfrac) and minimum number of psudo-detections (value of --minnumfalse).

-d FLT

Quantile limit to “grow” the final detections. As discussed in the previous options, after applying the initial quantile threshold, layers of pixels are carved off the objects to identify true signal. With this step you can return those low surface brightness layers that were carved off back to the detections. To disable growth, set the value of this option to 1.

The process is as follows: after the true detections are found, all the non-detected pixels above this quantile will be put in a list and used to “grow” the true detections (seeds of the growth). Like all quantile thresholds, this threshold is defined and applied to the convolved dataset. Afterwards, the dataset is dilated once (with minimum connectivity) to connect very thin regions on the boundary: imagine building a dam at the point rivers spill into an open sea/ocean. Finally, all holes are filled. In the geography metaphor, holes can be seen as the closed (by the dams) rivers and lakes, so this process is like turning the water in all such rivers and lakes into soil. See --detgrowmaxholesize for configuring the hole filling.


The maximum hole size to fill during the final expansion of the true detections as described in --detgrowquant. This is necessary when the input contains many smaller objects and can be used to avoid marking blank sky regions as detections.

For example multiple galaxies can be positioned such that they surround an empty region of sky. If all the holes are filled, the Sky region in between them will be taken as a detection which is not desired. To avoid such cases, the integer given to this option must be smaller than the hole between such objects. However, we should caution that unless the “hole” is very large, the combined faint wings of the galaxies might actually be present in between them, so be very careful in not filling such holes.

On the other hand, if you have a very large (and extended) galaxy, the diffuse wings of the galaxy may create very large holes over the detections. In such cases, a large enough value to this option will cause all such holes to be detected as part of the large galaxy and thus help in detecting it to extremely low surface brightness limits. Therefore, especially when large and extended objects are present in the image, it is recommended to give this option (very) large values. For one real-world example, see Detecting large extended targets.


After dilation, if the signal-to-noise ratio of a detection is less than the derived pseudo-detection S/N limit, that detection will be discarded. In an ideal/clean noise, a true detection’s S/N should be larger than its constituent pseudo-detections because its area is larger and it also covers more signal. However, on a false detections (especially at lower --snquant values), the increase in size can cause a decrease in S/N below that threshold.

This will improve purity and not change completeness (a true detection will not be discarded). Because a true detection has flux in its vicinity and dilation will catch more of that flux and increase the S/N. So on a true detection, the final S/N cannot be less than pseudo-detections.

However, in many real images bad processing creates artifacts that cannot be accurately removed by the Sky subtraction. In such cases, this option will decrease the completeness (will artificially discard true detections). So this feature is not default and should to be explicitly called when you know the noise is clean.


Every step of the detection process will be added as an extension to a file with the suffix _det.fits. Going through each would just be a repeat of the explanations above and also of those in Akhlaghi and Ichikawa (2015). The extension label should be sufficient to recognize which step you are observing. Viewing all the steps can be the best guide in choosing the best set of parameters. With this option, NoiseChisel will abort as soon as a snapshot of all the detection process is saved. This behavior can be disabled with --continueaftercheck.


Check the derivation of the final sky and its standard deviation values on the mesh grid. With this option, NoiseChisel will abort as soon as the sky value is estimated over the image (on each tile). This behavior can be disabled with --continueaftercheck. By default the output will have the same pixel size as the input, but with the --oneelempertile option, only one pixel will be used for each tile (see Processing options).

Previous: , Up: Invoking astnoisechisel   [Contents][Index] NoiseChisel output

NoiseChisel’s output is a multi-extension FITS file. The main extension/dataset is a (binary) detection map. It has the same size as the input but with only two possible values for all pixels: 0 (for pixels identified as noise) and 1 (for those identified as signal/detections). The detection map is followed by a Sky and Sky standard deviation dataset (which are calculated from the binary image). By default (when --rawoutput isn’t called), NoiseChisel will also subtract the Sky value from the input and save the sky-subtracted input as the first extension in the output.

The name of the output file can be set by giving a value to --output (this is a common option between all programs and is therefore discussed in Input/Output options). If --output isn’t used, the input name will be suffixed with _detected.fits and used as output, see Automatic output. If any of the options starting with --check* are given, NoiseChisel won’t complete and will abort as soon as the respective check images are created. For more information on the different check images, see the description for the --check* options in Detection options (this can be disabled with --continueaftercheck).

The last two extensions of the output are the Sky and its Standard deviation, see Sky value for a complete explanation. They are calculated on the tile grid that you defined for NoiseChisel. By default these datasets will have the same size as the input, but with all the pixels in one tile given one value. To be more space-efficient (keep only one pixel per tile), you can use the --oneelempertile option, see Tessellation.

To inspect any of NoiseChisel’s output files, assuming you use SAO DS9, you can configure your Graphic User Interface (GUI) to open NoiseChisel’s output as a multi-extension data cube. This will allow you to flip through the different extensions and visually inspect the results. This process has been described for the GNOME GUI (most common GUI in GNU/Linux operating systems) in Viewing multiextension FITS images.

NoiseChisel’s output configuration options are described in detail below.


Continue NoiseChisel after any of the options starting with --check (see Detection options. NoiseChisel involves many steps and as a result, there are many checks, allowing you to inspect the status of the processing. The results of each step affect the next steps of processing. Therefore, when you want to check the status of the processing at one step, the time spent to complete NoiseChisel is just wasted/distracting time.

To encourage easier experimentation with the option values, when you use any of the NoiseChisel options that start with --check, NoiseChisel will abort once its desired extensions have been written. With --continueaftercheck option, you can disable this behavior and ask NoiseChisel to continue with the rest of the processing, even after the requested check files are complete.


Don’t set the input’s blank pixels to blank in the output Sky and Sky standard deviation datasets. This is only applicable when the output has the same size as the input, in other words, when --oneelempertile isn’t called. By default, blank values in the input (commonly on the edges which are outside the survey/field area) will be set to blank in the output Sky and Sky standard deviation also.


Run a connected-components algorithm on the finally detected pixels to identify which pixels are connected to which. By default the main output is a binary dataset with only two values: 0 (for noise) and 1 (for signal/detections). See NoiseChisel output for more.

The purpose of NoiseChisel is to detect targets that are extended and diffuse, with outer parts that sink into the noise very gradually (galaxies and stars for example). Since NoiseChisel digs down to extremely low surface brightness values, many such targets will commonly be detected together as a single large body of connected pixels.

To properly separate connected objects, sophisticated segmentation methods are commonly necessary on NoiseChisel’s output. Gnuastro has the dedicated Segment program for this job. Since input images are commonly large and can take a significant volume, the extra volume necessary to store the labels of the connected components in the detection map (which will be created with this --label option, in 32-bit signed integer type) can thus be a major waste of space. Since the default output is just a binary dataset, an 8-bit unsigned dataset is enough.

The binary output will also encourage users to segment the result separately prior to doing higher-level analysis. As an alternative to --label, if you have the binary detection image, you can use the connected-components operator in Gnuastro’s Arithmetic program to identify regions that are connected with each other. For example with this command (assuming NoiseChisel’s output is called nc.fits):

$ astarithmetic nc.fits connected-components -hDETECTIONS

Don’t include the Sky-subtracted input image as the first extension of the output. By default, the Sky-subtracted input is put in the first extension of the output. The next extensions are NoiseChisel’s main outputs described above.

The extra Sky-subtracted input can be convenient in checking NoiseChisel’s output and comparing the detection map with the input: visually see if everything you expected is detected (reasonable completeness) and that you don’t have too many false detections (reasonable purity). This visual inspection is simplified if you use SAO DS9 to view NoiseChisel’s output as a multi-extension data-cube, see Viewing multiextension FITS images.

When you are satisfied with your NoiseChisel configuration (therefore you don’t need to check on every run), or you want to archive/transfer the outputs, or the datasets become large, or you are running NoiseChisel as part of a pipeline, this Sky-subtracted input image can be a significant burden (take up a large volume). The fact that the input is also noisy, makes it hard to compress it efficiently.

In such cases, this --rawoutput can be used to avoid the extra sky-subtracted input in the output. It is always possible to easily produce the Sky-subtracted dataset from the input (assuming it is in extension 1 of in.fits) and the SKY extension of NoiseChisel’s output (let’s call it nc.fits) with a command like below (assuming NoiseChisel wasn’t run with --oneelempertile, see Tessellation):

$ astarithmetic in.fits nc.fits - -h1 -hSKY

Save space: with the --rawoutput and --oneelempertile, NoiseChisel’s output will only be one binary detection map and two much smaller arrays with one value per tile. Since none of these have noise they can be compressed very effectively (without any loss of data) with exceptionally high compression ratios. This makes it easy to archive, or transfer, NoiseChisel’s output even on huge datasets. To compress it with the most efficient method (take up less volume), run the following command:

$ gzip --best noisechisel_output.fits

The resulting .fits.gz file can then be fed into any of Gnuastro’s programs directly, or viewed in viewers like SAO DS9, without having to decompress it separately (they will just take a little longer, because they have to internally decompress it before starting).

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7.3 Segment

Once signal is separated from noise (for example with NoiseChisel), you have a binary dataset: each pixel is either signal (1) or noise (0). Signal (for example every galaxy in your image) has been “detected”, but all detections have a label of 1. Therefore while we know which pixels contain signal, we still can’t find out how many galaxies they contain or which detected pixels correspond to which galaxy. At the lowest (most generic) level, detection is a kind of segmentation (segmenting the the whole dataset into signal and noise, see NoiseChisel). Here, we’ll define segmentation only on signal: to separate and find sub-structure within the detections.

If the targets are clearly separated, or their detected regions aren’t touching, a simple connected components128 algorithm (very basic segmentation) is enough to separate the regions that are touching/connected. This is such a basic and simple form of segmentation that Gnuastro’s Arithmetic program has an operator for it: see connected-components in Arithmetic operators. Assuming the binary dataset is called binary.fits, you can use it with a command like this:

$ astarithmetic binary.fits 2 connected-components

You can even do a very basic detection (a threshold, say at value 100) and segmentation in Arithmetic with a single command like below:

$ astarithmetic in.fits 100 gt 2 connected-components

However, in most astronomical situations our targets are not nicely separated or have a sharp boundary/edge (for a threshold to suffice): they touch (for example merging galaxies), or are simply in the same line-of-sight (which is much more common). This causes their images to overlap.

In particular, when you do your detection with NoiseChisel, you will detect signal to very low surface brightness limits: deep into the faint wings of galaxies or bright stars (which can extend very far and irregularly from their center). Therefore, it often happens that several galaxies are detected as one large detection. Since they are touching, a simple connected components algorithm will not suffice. It is therefore necessary to do a more sophisticated segmentation and break up the detected pixels (even those that are touching) into multiple target objects as accurately as possible.

Segment will use a detection map and its corresponding dataset to find sub-structure over the detected areas and use them for its segmentation. Until Gnuastro version 0.6 (released in 2018), Segment was part of NoiseChisel. Therefore, similar to NoiseChisel, the best place to start reading about Segment and understanding what it does (with many illustrative figures) is Section 3.2 of Akhlaghi and Ichikawa [2015].

As a summary, Segment first finds true clumps over the detections. Clumps are associated with local maxima/minima129 and extend over the neighboring pixels until they reach a local minimum/maximum (river/watershed). By default, Segment will use the distribution of clump signal-to-noise ratios over the undetected regions as reference to find “true” clumps over the detections. Using the undetected regions can be disabled by directly giving a signal-to-noise ratio to --clumpsnthresh.

The true clumps are then grown to a certain threshold over the detections. Based on the strength of the connections (rivers/watersheds) between the grown clumps, they are considered parts of one object or as separate objects. See Section 3.2 of Akhlaghi and Ichikawa [2015] (link above) for more. Segment’s main output are thus two labeled datasets: 1) clumps, and 2) objects. See Segment output for more.

To start learning about Segment, especially in relation to detection (NoiseChisel) and measurement (MakeCatalog), the recommended references are Akhlaghi and Ichikawa [2015] and Akhlaghi [2016].

Those papers cannot be updated any more but the software will evolve. For example Segment became a separate program (from NoiseChisel) in 2018 (after those papers were published). Therefore this book is the definitive reference. To help in the transition from those papers to the software you are using, see Segment changes after publication. Finally, in Invoking Segment, we’ll discuss Segment’s inputs, outputs and configuration options.

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7.3.1 Segment changes after publication

Segment’s main algorithm and working strategy were initially defined and introduced in Section 3.2 of Akhlaghi and Ichikawa [2015]. Prior to Gnuastro version 0.6 (released 2018), one program (NoiseChisel) was in charge of detection and segmentation. to increase creativity and modularity, NoiseChisel’s segmentation features were spun-off into a separate program (Segment). It is strongly recommended to read that paper for a good understanding of what Segment does, how it relates to detection, and how each parameter influences the output. That paper has a large number of figures showing every step on multiple mock and real examples.

However, the paper cannot be updated anymore, but Segment has evolved (and will continue to do so): better algorithms or steps have been (and will be) found. This book is thus the final and definitive guide to Segment. The aim of this section is to make the transition from the paper to your installed version, as smooth as possible through the list below. For a more detailed list of changes in previous Gnuastro releases/versions, please follow the NEWS file130.

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7.3.2 Invoking Segment

Segment will identify substructure within the detected regions of an input image. Segment’s output labels can be directly used for measurements (for example with MakeCatalog). The executable name is astsegment with the following general template

$ astsegment [OPTION ...] InputImage.fits

One line examples:

## Segment NoiseChisel's detected regions.
$ astsegment default-noisechisel-output.fits

## Use a hand-input S/N value for keeping true clumps
## (avoid finding the S/N using the undetected regions).
$ astsegment nc-out.fits --clumpsnthresh=10

## Inspect all the segmentation steps after changing a parameter.
$ astsegment input.fits --snquant=0.9 --checksegmentaion

## Use the fixed value of 0.01 for the input's Sky standard deviation
## (in the units of the input), and assume all the pixels are a
## detection (for example a large structure extending over the whole
## image), and only keep clumps with S/N>10 as true clumps.
$ astsegment in.fits --std=0.01 --detection=all --clumpsnthresh=10

If Segment is to do processing (for example you don’t want to get help, or see the values of each option), at least one input dataset is necessary along with detection and error information, either as separate datasets (per-pixel) or fixed values, see Segment input. Segment shares a large set of common operations with other Gnuastro programs, mainly regarding input/output, general processing steps, and general operating modes. To help in a unified experience between all of Gnuastro’s programs, these common operations have the same names and defined in Common options.

As in all Gnuastro programs, options can also be given to Segment in configuration files. For a thorough description of Gnuastro’s configuration file parsing, please see Configuration files. All of Segment’s options with a short description are also always available on the command-line with the --help option, see Getting help. To inspect the option values without actually running Segment, append your command with --printparams (or -P).

To help in easy navigation between Segment’s options, they are separately discussed in the three sub-sections below: Segment input discusses how you can customize the inputs to Segment. Segmentation options is devoted to options specific to the high-level segmentation process. Finally, in Segment output, we’ll discuss options that affect Segment’s output.

Next: , Previous: , Up: Invoking astsegment   [Contents][Index] Segment input

Besides the input dataset (for example astronomical image), Segment also needs to know the Sky standard deviation and the regions of the dataset that it should segment. The values dataset is assumed to be Sky subtracted by default. If it isn’t, you can ask Segment to subtract the Sky internally by calling --sky. For the rest of this discussion, we’ll assume it is already sky subtracted.

The Sky and its standard deviation can be a single value (to be used for the whole dataset) or a separate dataset (for a separate value per pixel). If a dataset is used for the Sky and its standard deviation, they must either be the size of the input image, or have a single value per tile (generated with --oneelempertile, see Processing options and Tessellation).

The detected regions/pixels can be specified as a detection map (for example see NoiseChisel output). If --detection=all, Segment won’t read any detection map and assume the whole input is a single detection. For example when the dataset is fully covered by a large nearby galaxy/globular cluster.

When dataset are to be used for any of the inputs, Segment will assume they are multiple extensions of a single file by default (when --std or --detection aren’t called). For example NoiseChisel’s default output NoiseChisel output. When the Sky-subtracted values are in one file, and the detection and Sky standard deviation are in another, you just need to use --detection: in the absence of --std, Segment will look for both the detection labels and Sky standard deviation in the file given to --detection. Ultimately, if all three are in separate files, you need to call both --detection and --std.

The extensions of the three mandatory inputs can be specified with --hdu, --dhdu, and --stdhdu. For a full discussion on what to give to these options, see the description of --hdu in Input/Output options. To see their default values (along with all the other options), run Segment with the --printparams (or -P) option. Just recall that in the absence of --detection and --std, all three are assumed to be in the same file. If you only want to see Segment’s default values for HDUs on your system, run this command:

$ astsegment -P | grep hdu

By default Segment will convolve the input with a kernel to improve the signal-to-noise ratio of true peaks. If you already have the convolved input dataset, you can pass it directly to Segment for faster processing (using the --convolved and --chdu options). Just don’t forget that the convolved image must also be Sky-subtracted before calling Segment. If a value/file is given to --sky, the convolved values will also be Sky subtracted internally. Alternatively, if you prefer to give a kernel (with --kernel and --khdu), Segment can do the convolution internally. To disable convolution, use --kernel=none.


The Sky value(s) to subtract from the input. This option can either be given a constant number or a file name containing a dataset (multiple values, per pixel or per tile). By default, Segment will assume the input dataset is Sky subtracted, so this option is not mandatory.

If the value can’t be read as a number, it is assumed to be a file name. When the value is a file, the extension can be specified with --skyhdu. When its not a single number, the given dataset must either have the same size as the output or the same size as the tessellation (so there is one pixel per tile, see Tessellation).

When this option is given, its value(s) will be subtracted from the input and the (optional) convolved dataset (given to --convolved) prior to starting the segmentation process.


The HDU/extension containing the Sky values. This is mandatory when the value given to --sky is not a number. Please see the description of --hdu in Input/Output options for the different ways you can identify a special extension.


The Sky standard deviation value(s) corresponding to the input. The value can either be a constant number or a file name containing a dataset (multiple values, per pixel or per tile). The Sky standard deviation is mandatory for Segment to operate.

If the value can’t be read as a number, it is assumed to be a file name. When the value is a file, the extension can be specified with --skyhdu. When its not a single number, the given dataset must either have the same size as the output or the same size as the tessellation (so there is one pixel per tile, see Tessellation).

When this option is not called, Segment will assume the standard deviation is a dataset and in a HDU/extension (--stdhdu) of another one of the input file(s). If a file is given to --detection, it will assume that file contains the standard deviation dataset, otherwise, it will look into input filename (the main argument, without any option).


The HDU/extension containing the Sky standard deviation values, when the value given to --std is a file name. Please see the description of --hdu in Input/Output options for the different ways you can identify a special extension.


The input Sky standard deviation value/dataset is actually variance. When this option is called, the square root of input Sky standard deviation (see --std) is used internally, not its raw value(s).

-d STR

Detection map to use for segmentation. If given a value of all, Segment will assume the whole dataset must be segmented, see below. If a detection map is given, the extension can be specified with --dhdu. If not given, Segment will assume the desired HDU/extension is in the main input argument (input file specified with no option).

The final segmentation (clumps or objects) will only be over the non-zero pixels of this detection map. The dataset must have the same size as the input image. Only datasets with an integer type are acceptable for the labeled image, see Numeric data types. If your detection map only has integer values, but it is stored in a floating point container, you can use Gnuastro’s Arithmetic program (see Arithmetic) to convert it to an integer container, like the example below:

$ astarithmetic float.fits int32 --output=int.fits

It may happen that the whole input dataset is covered by signal, for example when working on parts of the Andromeda galaxy, or nearby globular clusters (that cover the whole field of view). In such cases, segmentation is necessary over the complete dataset, not just specific regions (detections). By default Segment will first use the undetected regions as a reference to find the proper signal-to-noise ratio of “true” clumps (give a purity level specified with --snquant). Therefore, in such scenarios you also need to manually give a “true” clump signal-to-noise ratio with the --clumpsnthresh option to disable looking into the undetected regions, see Segmentation options. In such cases, is possible to make a detection map that only has the value 1 for all pixels (for example using Arithmetic), but for convenience, you can also use --detection=all.


The HDU/extension containing the detection map given to --detection. Please see the description of --hdu in Input/Output options for the different ways you can identify a special extension.

-k STR

The kernel used to convolve the input image. The usage of this option is identical to NoiseChisel’s --kernel option (NoiseChisel input). Please see the descriptions there for more. To disable convolution, you can give it a value of none.


The HDU/extension containing the kernel used for convolution. For acceptable values, please see the description of --hdu in Input/Output options.


The convolved image to avoid internal convolution by Segment. The usage of this option is identical to NoiseChisel’s --convolved option. Please see NoiseChisel input for a thorough discussion of the usefulness and best practices of using this option.

If you want to use the same convolution kernel for detection (with NoiseChisel) and segmentation, with this option, you can use the same convolved image (that is also available in NoiseChisel) and avoid two convolutions. However, just be careful to use the input to NoiseChisel as the input to Segment also, then use the --sky and --std to specify the Sky and its standard deviation (from NoiseChisel’s output). Recall that when NoiseChisel is not called with --rawoutput, the first extension of NoiseChisel’s output is the Sky-subtracted input (see NoiseChisel output). So if you use the same convolved image that you fed to NoiseChisel, but use NoiseChisel’s output with Segment’s --convolved, then the convolved image won’t be Sky subtracted.


The HDU/extension containing the convolved image (given to --convolved). For acceptable values, please see the description of --hdu in Input/Output options.


The size of the large tiles to use for identifying the clump S/N threshold over the undetected regions. The usage of this option is identical to NoiseChisel’s --largetilesize option (NoiseChisel input). Please see the descriptions there for more.

The undetected regions can be a significant fraction of the dataset and finding clumps requires sorting of the desired regions, which can be slow. To speed up the processing, Segment finds clumps in the undetected regions over separate large tiles. This allows it to have to sort a much smaller set of pixels and also to treat them independently and in parallel. Both these issues greatly speed it up. Just be sure to not decrease the large tile sizes too much (less than 100 pixels in each dimension). It is important for them to be much larger than the clumps.

Next: , Previous: , Up: Invoking astsegment   [Contents][Index] Segmentation options

The options below can be used to configure every step of the segmentation process in the Segment program. For a more complete explanation (with figures to demonstrate each step), please see Section 3.2 of Akhlaghi and Ichikawa [2015], and also Segment. By default, Segment will follow the procedure described in the paper to find the S/N threshold based on the noise properties. This can be disabled by directly giving a trustable signal-to-noise ratio to the --clumpsnthresh option.

Recall that you can always see the full list of Gnuastro’s options with the --help (see Getting help), or --printparams (or -P) to see their values (see Operating mode options).


Minimum fraction (value between 0 and 1) of Sky (undetected) areas in a large tile. Only (large) tiles with a fraction of undetected pixels (Sky) greater than this value will be used for finding clumps. The clumps found in the undetected areas will be used to estimate a S/N threshold for true clumps. Therefore this is an important option (to decrease) in crowded fields. Operationally, this is almost identical to NoiseChisel’s --minskyfrac option (Detection options). Please see the descriptions there for more.


Build the clumps based on the local minima, not maxima. By default, clumps are built starting from local maxima (see Figure 8 of Akhlaghi and Ichikawa [2015]). Therefore, this option can be useful when you are searching for true local minima (for example absorption features).

-m INT

The minimum area which a clump in the undetected regions should have in order to be considered in the clump Signal to noise ratio measurement. If this size is set to a small value, the Signal to noise ratio of false clumps will not be accurately found. It is recommended that this value be larger than the value to NoiseChisel’s --snminarea. Because the clumps are found on the convolved (smoothed) image while the psudo-detections are found on the input image. You can use --checksn and --checksegmentation to see if your chosen value is reasonable or not.


Save the S/N values of the clumps into two files ending with _clumpsn_sky.XXX and _clumpsn_det.XXX. The .XXX is determined from the --tableformat option (see Input/Output options, for example .txt or .fits). You can use these to inspect the S/N values and their distribution (in combination with the --checksegmentation option to see where the clumps are). You can use Gnuastro’s Statistics to make a histogram of the distribution (ready for plotting in a text file, or a crude ASCII-art demonstration on the command-line).

With this option, NoiseChisel will abort as soon as the two tables are created. This allows you to inspect the steps leading to the final S/N quantile threshold, this behavior can be disabled with --continueaftercheck.


The minimum number of clumps over undetected (Sky) regions to identify the requested Signal-to-Noise ratio threshold. Operationally, this is almost identical to NoiseChisel’s --minnumfalse option (Detection options). Please see the descriptions there for more.

-c FLT

The quantile of the signal-to-noise ratio distribution of clumps in undetected regions, used to define true clumps. After identifying all the usable clumps in the undetected regions of the dataset, the given quantile of their signal-to-noise ratios is used to define a the signal-to-noise ratio of a “true” clump. Effectively, this can be seen as an inverse p-value measure. See Figure 9 and Section 3.2.1 of Akhlaghi and Ichikawa [2015] for a complete explanation. The full distribution of clump signal-to-noise ratios over the undetected areas can be saved into a table with --checksn option and visually inspected with --checksegmentation.


Keep a clump whose maximum (minimum if --minima is called) flux is 8-connected to a river pixel. By default such clumps over detections are considered to be noise and are removed irrespective of their brightness (see Flux Brightness and magnitude). Over large profiles, that sink into the noise very slowly, noise can cause part of the profile (which was flat without noise) to become a very large and with a very high Signal to noise ratio. In such cases, the pixel with the maximum flux in the clump will be immediately touching a river pixel.

-s FLT

The signal-to-noise threshold for true clumps. If this option is given, then the segmentation options above will be ignored and the given value will be directly used to identify true clumps over the detections. This can be useful if you have a large dataset with similar noise properties. You can find a robust signal-to-noise ratio based on a (sufficiently large) smaller portion of the dataset. Afterwards, with this option, you can speed up the processing on the whole dataset. Other scenarios where this option may be useful is when, the image might not contain enough/any Sky regions.


Threshold (multiple of the sky standard deviation added with the sky) to stop growing true clumps. Once true clumps are found, they are set as the basis to segment the detected region. They are grown until the threshold specified by this option.

-y INT

The minimum length of a river between two grown clumps for it to be considered in signal-to-noise ratio estimations. Similar to --snminarea, if the length of the river is too short, the signal-to-noise ratio can be noisy and unreliable. Any existing rivers shorter than this length will be considered as non-existent, independent of their Signal to noise ratio. The clumps are grown on the input image, therefore this value can be smaller than the value given to --snminarea. Recall that the clumps were defined on the convolved image so --snminarea should be larger.


The maximum Signal to noise ratio of the rivers between two grown clumps in order to consider them as separate ‘objects’. If the Signal to noise ratio of the river between two grown clumps is larger than this value, they are defined to be part of one ‘object’. Note that the physical reality of these ‘objects’ can never be established with one image, or even multiple images from one broad-band filter. Any method we devise to define ‘object’s over a detected region is ultimately subjective.

Two very distant galaxies or satellites in one halo might lie in the same line of sight and be detected as clumps on one detection. On the other hand, the connection (through a spiral arm or tidal tail for example) between two parts of one galaxy might have such a low surface brightness that they are broken up into multiple detections or objects. In fact if you have noticed, exactly for this purpose, this is the only Signal to noise ratio that the user gives into NoiseChisel. The ‘true’ detections and clumps can be objectively identified from the noise characteristics of the image, so you don’t have to give any hand input Signal to noise ratio.


A file with the suffix _seg.fits will be created. This file keeps all the relevant steps in finding true clumps and segmenting the detections into multiple objects in various extensions. Having read the paper or the steps above. Examining this file can be an excellent guide in choosing the best set of parameters. Note that calling this function will significantly slow NoiseChisel. In verbose mode (without the --quiet option, see Operating mode options) the important steps (along with their extension names) will also be reported.

With this option, NoiseChisel will abort as soon as the two tables are created. This behavior can be disabled with --continueaftercheck.

Previous: , Up: Invoking astsegment   [Contents][Index] Segment output

The main output of Segment are two label datasets (with integer types, separating the dataset’s elements into different classes) with HDU/extension names of CLUMPS and OBJECTS. For a complete definition of clumps and objects, please see Section 3.2 of Akhlaghi and Ichikawa [2015] and Segmentation options.

The clumps are “true” local maxima (minima if --minima is called) and their surrounding pixels until a local minimum/maximum (caused by noise fluctuations, or another “true” clump). Therefore it may happen that some of the input detections aren’t covered by clumps at all (very diffuse objects without any strong peak), while some objects may contain many clumps. Even in those that have clumps, there will be regions that are too diffuse. The diffuse regions (within the input detected regions) are given a negative label (-1) to help you separate them from the undetected regions (with a value of zero).

Each clump is labeled with respect to its host object. Therefore, if an object has three clumps for example, the clumps within it have labels 1, 2 and 3. As a result, if an initial detected region has multiple objects, each with a single clump, all the clumps will have a label of 1. The total number of clumps in the dataset is stored in the NCLUMPS keyword of the CLUMPS extension and printed in the verbose output of Segment (when --quiet is not called).

The OBJECTS extension of the output will give a positive counter/label to every detected pixel in the input. As described in Akhlaghi and Ichikawa [2015], the true clumps are grown until a certain threshold. If the grown clumps touch other clumps and the connection is strong enough, they are considered part of the same object. Once objects (grown clumps) are identified, they are grown to cover the whole detected area.

By default, besides the CLUMPS and OBJECTS extensions, Segment’s output will also contain the (technically redundant) input dataset and the sky standard deviation dataset (if it wasn’t a constant number). This can help in visually inspecting the result when viewing the images as a “Multi-extension data cube” in SAO DS9 for example (see Viewing multiextension FITS images). You can simply flip through the extensions and see the same region of the image and its corresponding clumps/object labels. It also makes it easy to feed the output (as one file) into MakeCatalog when you intend to make a catalog afterwards (see MakeCatalog. To remove these redundant extensions from the output (for example when designing a pipeline), you can use --rawoutput.

The OBJECTS and CLUMPS extensions can be used as input into MakeCatalog to generate a catalog for higher-level analysis. If you want to treat each clump separately, you can give a very large value (or even a NaN, which will always fail) to the --gthresh option (for example --gthresh=1e10 or --gthresh=nan), see Segmentation options. The options to configure the output of Segment are listed below:


Don’t abort Segment after producing the check image(s). The usage of this option is identical to NoiseChisel’s --continueaftercheck option (NoiseChisel input). Please see the descriptions there for more.


Abort Segment after finding true clumps and don’t continue with finding options. Therefore, no OBJECTS extension will be present in the output. Each true clump in CLUMPS will get a unique label, but diffuse regions will still have a negative value.

To make a catalog of the clumps, the input detection map (where all the labels are one) can be fed into MakeCatalog along with the input detection map to Segment (that only had a value of 1 for all detected pixels) with --clumpscat. In this way, MakeCatalog will assume all the clumps belong to a single “object”.


In the output CLUMPS extension, store the grown clumps. If a detected region contains no clumps or only one clump, then it will be fully given a label of 1 (no negative valued pixels).


Only write the CLUMPS and OBJECTS datasets in the output file. Without this option (by default), the first and last extensions of the output will the Sky-subtracted input dataset and the Sky standard deviation dataset (if it wasn’t a number). When the datasets are small, these redundant extensions can make it convenient to inspect the results visually or feed the output to MakeCatalog for measurements. Ultimately both the input and Sky standard deviation datasets are redundant (you had them before running Segment). When the inputs are large/numerous, these extra dataset can be a burden.

Save space: with the --rawoutput, Segment’s output will only be two labeled datasets (only containing integers). Since they have no noise, such datasets can be compressed very effectively (without any loss of data) with exceptionally high compression ratios. You can use the following command to compress it with the best ratio:

$ gzip --best segment_output.fits

The resulting .fits.gz file can then be fed into any of Gnuastro’s programs directly, without having to decompress it separately (it will just take them a little longer, because they have to decompress it internally before use).

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7.4 MakeCatalog

At the lowest level, a dataset (for example an image) is just a collection of values, placed after each other in any number of dimensions (for example an image is a 2D dataset). Each data-element (pixel) just has two properties: its position (relative to the rest) and its value. In higher-level analysis, an entire dataset (an image for example) is rarely treated as a singular entity133. You usually want to know/measure the properties of the (separate) scientifically interesting targets that are embedded in it. For example the magnitudes, positions and elliptical properties of the galaxies that are in the image.

MakeCatalog is Gnuastro’s program for localized measurements over a dataset. The role of MakeCatalog in a scientific analysis and the benefits of this model of data analysis (were detection/segmentation is separated from measurement) is discussed in Akhlaghi [2016]. We strongly recommend reading this short paper for a better understanding of this methodology thus effective usage of MakeCatalog, in combination with the other Gnuastro’s programs. However, that paper cannot undergo any more change, so this manual is the definitive guide.

It is important to define your regions of interest before running MakeCatalog. MakeCatalog is specialized in doing measurements accurately and efficiently. Therefore MakeCatalog will not do detection, segmentation, or defining apertures on requested positions in your dataset. Following Gnuastro’s modularity principle, There are separate and highly specialized and customizable programs in Gnuastro for these other jobs:

These programs will/can return labeled dataset(s) to be fed into MakeCatalog. The labeled dataset must have the same size/dimensions as the input, but only with integer valued pixels that have the label/counter for the feature the pixel belongs to.

These labels are then directly used to make the necessary measurements. For example the flux weighted average position of all the pixels with a label of 42 will be considered as the central position134 of the 42nd row of the output catalog. Similarly, the sum of all these pixels will be the 42nd row in the brightness column and etc. Pixels with labels equal to or smaller than zero will be ignored by MakeCatalog. In other words, the number of rows in MakeCatalog’s output is already known before running it.

Before getting into the details of running MakeCatalog (in Invoking MakeCatalog, we’ll start with a discussion on the basics of its approach to separating detection from measurements in Detection and catalog production. A very important factor in any measurement is understanding its validity range, or limits. Therefore in Quantifying measurement limits, we’ll discuss how to estimate the reliability of the detection and basic measurements. This section will continue with a derivation of elliptical parameters from the labeled datasets in Measuring elliptical parameters. For those who feel MakeCatalog’s existing measurements/columns aren’t enough and would like to add further measurements, in Adding new columns to MakeCatalog, a checklist of steps is provided for readily adding your own new measurements/columns.

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7.4.1 Detection and catalog production

Most other common tools in low-level astronomical data-analysis (for example SExtractor135) merge the two processes of detection and measurement into one. Gnuastro’s modularized methodology to separating detection from measurements is therefore new to many experienced astronomers and deserves a short review here. Further discussion on the benefits of this methodology can be seen in Akhlaghi [2016].

The raw output of any of the programs mentioned in MakeCatalog can be directly fed into MakeCatalog to get to a fast catalog. This is good when no further customization is necessary and you want a fast/simple catalog. But the modular approach taken by Gnuastro has many benefits that will become clear as you get more experienced in astronomical data analysis and want to be more creative in using your valuable data for the exciting scientific project you are working on. In short the reasons for this modularity can be classified as below:

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7.4.2 Quantifying measurement limits

No measurement on a real dataset can be perfect: you can only reach a certain level/limit of accuracy. Therefore, a meaningful (scientific) analysis requires an understanding of these limits for the dataset and your analysis tools: different datasets have different noise properties and different detection methods (one method/algorith/software that is run with a different set of parameters is considered as a different detection method) will have different abilities to detect or measure certain kinds of signal (astronomical objects) and their properties in the dataset. Hence, quantifying the detection and measurement limitations with a particular dataset and analysis tool is the most crucial/critical aspect of any high-level analysis.

Here, we’ll review some of the most general limits that are important in any astronomical data analysis and how MakeCatalog makes it easy to find them. Depending on the higher-level analysis, there are more tests that must be done, but these are relatively low-level and usually necessary in most cases. In astronomy, it is common to use the magnitude (a unit-less scale) and physical units, see Flux Brightness and magnitude. Therefore the measurements discussed here are commonly used in units of magnitudes.

Surface brightness limit (of whole dataset)

As we make more observations on one region of the sky, and add the observations into one dataset, the signal and noise both increase. However, the signal increase much faster than the noise: assuming you add \(N\) datasets with equal exposure times, the signal will increases as a multiple of \(N\), while noise increases as \(\sqrt{N}\). Thus this increases the signal-to-noise ratio. Qualitatively, fainter (per pixel) parts of the objects/signal in the image will become more visible/detectable. The noise-level is known as the dataset’s surface brightness limit.

You can think of the noise as muddy water that is completely covering a flat ground136. The signal (or astronomical objects in this analogy) will be summits/hills that start from the flat sky level (under the muddy water) and can sometimes reach outside of the muddy water. Let’s assume that in your first observation the muddy water has just been stirred and you can’t see anything through it. As you wait and make more observations/exposures, the mud settles down and the depth of the transparent water increases, making the summits visible. As the depth of clear water increases, the parts of the hills with lower heights (parts with lower surface brightness) can be seen more clearly. In this analogy, height (from the ground) is surface brightness137 and the height of the muddy water is your surface brightness limit.

The outputs of NoiseChisel include the Sky standard deviation (\(\sigma\)) on every group of pixels (a mesh) that were calculated from the undetected pixels in each tile, see Tessellation and NoiseChisel output. Let’s take \(\sigma_m\) as the median \(\sigma\) over the successful meshes in the image (prior to interpolation or smoothing).

On different instruments, pixels have different physical sizes (for example in micro-meters, or spatial angle over the sky). Nevertheless, a pixel is our unit of data collection. In other words, while quantifying the noise, the physical or projected size of the pixels is irrelevant. We thus define the Surface brightness limit or depth, in units of magnitude/pixel, of a data-set, with zeropoint magnitude \(z\), with the \(n\)th multiple of \(\sigma_m\) as (see Flux Brightness and magnitude):

$$SB_{\rm Pixel}=-2.5\times\log_{10}{(n\sigma_m)}+z$$

As an example, the XDF survey covers part of the sky that the Hubble space telescope has observed the most (for 85 orbits) and is consequently very small (\(\sim4\) arcmin\(^2\)). On the other hand, the CANDELS survey, is one of the widest multi-color surveys covering several fields (about 720 arcmin\(^2\)) but its deepest fields have only 9 orbits observation. The depth of the XDF and CANDELS-deep surveys in the near infrared WFC3/F160W filter are respectively 34.40 and 32.45 magnitudes/pixel. In a single orbit image, this same field has a depth of 31.32. Recall that a larger magnitude corresponds to less brightness.

The low-level magnitude/pixel measurement above is only useful when all the datasets you want to use belong to one instrument (telescope and camera). However, you will often find yourself using datasets from various instruments with different pixel scales (projected pixel sizes). If we know the pixel scale, we can obtain a more easily comparable surface brightness limit in units of: magnitude/arcsec\(^2\). Let’s assume that the dataset has a zeropoint value of \(z\), and every pixel is \(p\) arcsec\(^2\) (so \(A/p\) is the number of pixels that cover an area of \(A\) arcsec\(^2\)). If the surface brightness is desired at the \(n\)th multiple of \(\sigma_m\), the following equation (in units of magnitudes per A arcsec\(^2\)) can be used:

$$SB_{\rm Projected}=-2.5\times\log_{10}{\left(n\sigma_m\sqrt{A\over p}\right)+z}$$

Note that this is just an extrapolation of the per-pixel measurement \(\sigma_m\). So it should be used with extreme care: for example the dataset must have an approximately flat depth or noise properties overall. A more accurate measure for each detection over the dataset is known as the upper-limit magnitude which actually uses random positioning of each detection’s area/footprint (see below). It doesn’t extrapolate and even accounts for correlated noise patterns in relation to that detection. Therefore, the upper-limit magnitude is a much better measure of your dataset’s surface brightness limit for each particular object.

MakeCatalog will calculate the input dataset’s \(SB_{\rm Pixel}\) and \(SB_{\rm Projected}\) and write them as comments/meta-data in the output catalog(s). Just note that \(SB_{\rm Projected}\) is only calculated if the input has World Coordinate System (WCS).

Completeness limit (of each detection)

As the surface brightness of the objects decreases, the ability to detect them will also decrease. An important statistic is thus the fraction of objects of similar morphology and brightness that will be identified with our detection algorithm/parameters in the given image. This fraction is known as completeness. For brighter objects, completeness is 1: all bright objects that might exist over the image will be detected. However, as we go to objects of lower overall surface brightness, we will fail to detect some, and gradually we are not able to detect anything any more. For a given profile, the magnitude where the completeness drops below a certain level (usually above \(90\%\)) is known as the completeness limit.

Another important parameter in measuring completeness is purity: the fraction of true detections to all true detections. In effect purity is the measure of contamination by false detections: the higher the purity, the lower the contamination. Completeness and purity are anti-correlated: if we can allow a large number of false detections (that we might be able to remove by other means), we can significantly increase the completeness limit.

One traditional way to measure the completeness and purity of a given sample is by embedding mock profiles in regions of the image with no detection. However in such a study we must be really careful to choose model profiles as similar to the target of interest as possible.

Magnitude measurement error (of each detection)

Any measurement has an error and this includes the derived magnitude for an object. Note that this value is only meaningful when the object’s magnitude is brighter than the upper-limit magnitude (see the next items in this list). As discussed in Flux Brightness and magnitude, the magnitude (\(M\)) of an object with brightness \(B\) and Zeropoint magnitude \(z\) can be written as:


Calculating the derivative with respect to \(B\), we get:

$${dM\over dB} = {-2.5\over {B\times ln(10)}}$$

From the Tailor series (\(\Delta{M}=dM/dB\times\Delta{B}\)), we can write:

$$\Delta{M} = \left|{-2.5\over ln(10)}\right|\times{\Delta{B}\over{B}}$$

But, \(\Delta{B}/B\) is just the inverse of the Signal-to-noise ratio (\(S/N\)), so we can write the error in magnitude in terms of the signal-to-noise ratio:

$$\Delta{M} = {2.5\over{S/N\times ln(10)}} $$

MakeCatalog uses this relation to estimate the magnitude errors. The signal-to-noise ratio is calculated in different ways for clumps and objects (see Akhlaghi and Ichikawa [2015]), but this single equation can be used to estimate the measured magnitude error afterwards for any type of target.

Upper limit magnitude (of each detection)

Due to the noisy nature of data, it is possible to get arbitrarily low values for a faint object’s brightness (or arbitrarily high magnitudes). Given the scatter caused by the dataset’s noise, values fainter than a certain level are meaningless: another similar depth observation will give a radically different value.

For example, while the depth of the image is 32 magnitudes/pixel, a measurement that gives a magnitude of 36 for a \(\sim100\) pixel object is clearly unreliable. In another similar depth image, we might measure a magnitude of 30 for it, and yet another might give 33. Furthermore, due to the noise scatter so close to the depth of the data-set, the total brightness might actually get measured as a negative value, so no magnitude can be defined (recall that a magnitude is a base-10 logarithm). This problem usually becomes relevant when the detection labels were not derived from the values being measured (for example when you are estimating colors, see MakeCatalog).

Using such unreliable measurements will directly affect our analysis, so we must not use the raw measurements. But how can we know how reliable a measurement on a given dataset is?

When we confront such unreasonably faint magnitudes, there is one thing we can deduce: that if something actually exists here (possibly buried deep under the noise), it’s inherent magnitude is fainter than an upper limit magnitude. To find this upper limit magnitude, we place the object’s footprint (segmentation map) over random parts of the image where there are no detections, so we only have pure (possibly correlated) noise, along with undetected objects. Doing this a large number of times will give us a distribution of brightness values. The standard deviation (\(\sigma\)) of that distribution can be used to quantify the upper limit magnitude.

Traditionally, faint/small object photometry was done using fixed circular apertures (for example with a diameter of \(N\) arc-seconds). Hence, the upper limit was like the depth discussed above: one value for the whole image. The problem with this simplified approach is that the number of pixels in the aperture directly affects the final distribution and thus magnitude. Also the image correlated noise might actually create certain patters, so the shape of the object can also affect the final result result. Fortunately, with the much more advanced hardware and software of today, we can make customized segmentation maps for each object.

When requested, MakeCatalog will randomly place each target’s footprint over the dataset as described above and estimate the resulting distribution’s properties (like the upper limit magnitude). The procedure is fully configurable with the options in Upper-limit settings. If one value for the whole image is required, you can either use the surface brightness limit above or make a circular aperture and feed it into MakeCatalog to request an upper-limit magnitude for it138.

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7.4.3 Measuring elliptical parameters

The shape or morphology of a target is one of the most commonly desired parameters of a target. Here, we will review the derivation of the most basic/simple morphological parameters: the elliptical parameters for a set of labeled pixels. The elliptical parameters are: the (semi-)major axis, the (semi-)minor axis and the position angle along with the central position of the profile. The derivations below follow the SExtractor manual derivations with some added explanations for easier reading.

Let’s begin with one dimension for simplicity: Assume we have a set of \(N\) values \(B_i\) (keeping the spatial distribution of brightness for example), each at position \(x_i\). The simplest parameter we can define is the geometric center of the object (\(x_g\)) (ignoring the brightness values): \(x_g=(\sum_ix_i)/N\). Moments are defined to incorporate both the value (brightness) and position of the data. The first moment can be written as:

$$\overline{x}={\sum_iB_ix_i \over \sum_iB_i}$$

This is essentially the weighted (by \(B_i\)) mean position. The geometric center (\(x_g\), defined above) is a special case of this with all \(B_i=1\). The second moment is essentially the variance of the distribution:

$$\overline{x^2}\equiv{\sum_iB_i(x_i-\overline{x})^2 \over \sum_iB_i} = {\sum_iB_ix_i^2 \over \sum_iB_i} - 2\overline{x}{\sum_iB_ix_i\over\sum_iB_i} + \overline{x}^2 ={\sum_iB_ix_i^2 \over \sum_iB_i} - \overline{x}^2$$

The last step was done from the definition of \(\overline{x}\). Hence, the square root of \(\overline{x^2}\) is the spatial standard deviation (along the one-dimension) of this particular brightness distribution (\(B_i\)). Crudely (or qualitatively), you can think of its square root as the distance (from \(\overline{x}\)) which contains a specific amount of the flux (depending on the \(B_i\) distribution). Similar to the first moment, the geometric second moment can be found by setting all \(B_i=1\). So while the first moment quantified the position of the brightness distribution, the second moment quantifies how that brightness is dispersed about the first moment. In other words, it quantifies how “sharp” the object’s image is.

Before continuing to two dimensions and the derivation of the elliptical parameters, let’s pause for an important implementation technicality. You can ignore this paragraph and the next two if you don’t want to implement these concepts. The basic definition (first definition of \(\overline{x^2}\) above) can be used without any major problem. However, using this fraction requires two runs over the data: one run to find \(\overline{x}\) and another run to find \(\overline{x^2}\) from \(\overline{x}\), this can be slow. The advantage of the last fraction above, is that we can estimate both the first and second moments in one run (since the \(-\overline{x}^2\) term can easily be added later).

The logarithmic nature of floating point number digitization creates a complication however: suppose the object is located between pixels 10000 and 10020. Hence the target’s pixels are only distributed over 20 pixels (with a standard deviation \(<20\)), while the mean has a value of \(\sim10000\). The \(\sum_iB_i^2x_i^2\) will go to very very large values while the individual pixel differences will be orders of magnitude smaller. This will lower the accuracy of our calculation due to the limited accuracy of floating point operations. The variance only depends on the distance of each point from the mean, so we can shift all position by a constant/arbitrary \(K\) which is much closer to the mean: \(\overline{x-K}=\overline{x}-K\). Hence we can calculate the second order moment using:

$$\overline{x^2}={\sum_iB_i(x_i-K)^2 \over \sum_iB_i} - (\overline{x}-K)^2 $$

The closer \(K\) is to \(\overline{x}\), the better (the sums of squares will involve smaller numbers), as long as \(K\) is within the object limits (in the example above: \(10000\leq{K}\leq10020\)), the floating point error induced in our calculation will be negligible. For the most simplest implementation, MakeCatalog takes \(K\) to be the smallest position of the object in each dimension. Since \(K\) is arbitrary and an implementation/technical detail, we will ignore it for the remainder of this discussion.

In two dimensions, the mean and variances can be written as:

$$\overline{x}={\sum_iB_ix_i\over B_i}, \quad \overline{x^2}={\sum_iB_ix_i^2 \over \sum_iB_i} - \overline{x}^2$$ $$\overline{y}={\sum_iB_iy_i\over B_i}, \quad \overline{y^2}={\sum_iB_iy_i^2 \over \sum_iB_i} - \overline{y}^2$$ $$\quad\quad\quad\quad\quad\quad\quad\quad\quad \overline{xy}={\sum_iB_ix_iy_i \over \sum_iB_i} - \overline{x}\times\overline{y}$$

If an elliptical profile’s major axis exactly lies along the \(x\) axis, then \(\overline{x^2}\) will be directly proportional with the profile’s major axis, \(\overline{y^2}\) with its minor axis and \(\overline{xy}=0\). However, in reality we are not that lucky and (assuming galaxies can be parameterized as an ellipse) the major axis of galaxies can be in any direction on the image (in fact this is one of the core principles behind weak-lensing by shear estimation). So the purpose of the remainder of this section is to define a strategy to measure the position angle and axis ratio of some randomly positioned ellipses in an image, using the raw second moments that we have calculated above in our image coordinates.

Let’s assume we have rotated the galaxy by \(\theta\), the new second order moments are:

$$\overline{x_\theta^2} = \overline{x^2}\cos^2\theta + \overline{y^2}\sin^2\theta - 2\overline{xy}\cos\theta\sin\theta $$ $$\overline{y_\theta^2} = \overline{x^2}\sin^2\theta + \overline{y^2}\cos^2\theta + 2\overline{xy}\cos\theta\sin\theta$$ $$\overline{xy_\theta} = \overline{x^2}\cos\theta\sin\theta - \overline{y^2}\cos\theta\sin\theta + \overline{xy}(\cos^2\theta-\sin^2\theta)$$

The best \(\theta\) (\(\theta_0\), where major axis lies along the \(x_\theta\) axis) can be found by:

$$\left.{\partial \overline{x_\theta^2} \over \partial \theta}\right|_{\theta_0}=0$$ Taking the derivative, we get: $$2\cos\theta_0\sin\theta_0(\overline{y^2}-\overline{x^2}) + 2(\cos^2\theta_0-\sin^2\theta_0)\overline{xy}=0$$ When \(\overline{x^2}\neq\overline{y^2}\), we can write: $$\tan2\theta_0 = 2{\overline{xy} \over \overline{x^2}-\overline{y^2}}.$$

MakeCatalog uses the standard C math library’s atan2 function to estimate \(\theta_0\), which we define as the position angle of the ellipse. To recall, this is the angle of the major axis of the ellipse with the \(x\) axis. By definition, when the elliptical profile is rotated by \(\theta_0\), then \(\overline{xy_{\theta_0}}=0\), \(\overline{x_{\theta_0}^2}\) will be the extent of the maximum variance and \(\overline{y_{\theta_0}^2}\) the extent of the minimum variance (which are perpendicular for an ellipse). Replacing \(\theta_0\) in the equations above for \(\overline{x_\theta}\) and \(\overline{y_\theta}\), we can get the semi-major (\(A\)) and semi-minor (\(B\)) lengths:

$$A^2\equiv\overline{x_{\theta_0}^2}= {\overline{x^2} + \overline{y^2} \over 2} + \sqrt{\left({\overline{x^2}-\overline{y^2} \over 2}\right)^2 + \overline{xy}^2}$$

$$B^2\equiv\overline{y_{\theta_0}^2}= {\overline{x^2} + \overline{y^2} \over 2} - \sqrt{\left({\overline{x^2}-\overline{y^2} \over 2}\right)^2 + \overline{xy}^2}$$

As a summary, it is important to remember that the units of \(A\) and \(B\) are in pixels (the standard deviation of a positional distribution) and that they represent the spatial light distribution of the object in both image dimensions (rotated by \(\theta_0\)). When the object cannot be represented as an ellipse, this interpretation breaks down: \(\overline{xy_{\theta_0}}\neq0\) and \(\overline{y_{\theta_0}^2}\) will not be the direction of minimum variance.

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7.4.4 Adding new columns to MakeCatalog

MakeCatalog is designed to allow easy addition of different measurements over a labeled image (see Akhlaghi [2016]). A check-list style description of necessary steps to do that is described in this section. The common development characteristics of MakeCatalog and other Gnuastro programs is explained in Developing. We strongly encourage you to have a look at that chapter to greatly simplify your navigation in the code. After adding and testing your column, you are most welcome (and encouraged) to share it with us so we can add to the next release of Gnuastro for everyone else to also benefit from your efforts.

MakeCatalog will first pass over each label’s pixels two times and do necessary raw/internal calculations. Once the passes are done, it will use the raw information for filling the final catalog’s columns. In the first pass it will gather mainly object information and in the second run, it will mainly focus on the clumps, or any other measurement that needs an output from the first pass. These two passes are designed to be raw summations: no extra processing. This will allow parallel processing and simplicity/clarity. So if your new calculation, needs new raw information from the pixels, then you will need to also modify the respective mkcatalog_first_pass and mkcatalog_second_pass functions (both in bin/mkcatalog/mkcatalog.c) and define new raw table columns in main.h (hopefully the comments in the code are clear enough).

In all these different places, the final columns are sorted in the same order (same order as Invoking MakeCatalog). This allows a particular column/option to be easily found in all steps. Therefore in adding your new option, be sure to keep it in the same relative place in the list in all the separate places (it doesn’t necessarily have to be in the end), and near conceptually similar options.


The objectcols and clumpcols enumerated variables (enum) define the raw/internal calculation columns. If your new column requires new raw calculations, add a row to the respective list. If your calculation requires any other settings parameters, you should add a variable to the mkcatalogparams structure.


If the new column needs raw calculations (an entry was added in objectcols and clumpcols), specify which inputs it needs in ui_necessary_inputs, similar to the other options. Afterwards, if your column includes any particular settings (you needed to add a variable to the mkcatalogparams structure in main.h), you should do the sanity checks and preparations for it here.


The option_keys_enum associates a unique value for each option to MakeCatalog. The options that have a short option version, the single character short comment is used for the value. Those that don’t have a short option version, get a large integer automatically. You should add a variable here to identify your desired column.


This file specifies all the parameters for the GNU C library, Argp structure that is in charge of reading the user’s options. To define your new column, just copy an existing set of parameters and change the first, second and 5th values (the only ones that differ between all the columns), you should use the macro you defined in ui.h here.


This file contains the main definition and high-level calculation of your new column through the columns_define_alloc and columns_fill functions. In the first, you specify the basic information about the column: its name, units, comments, type (see