MMTF Astronomer's Guide
Tunable Filter Basic Principles
The Advantages of Tunable Filter Imaging

Fabry-Perot tunable filters operate on similar principles to narrow-band interference filters. In a single-plate interference filter, the interference occurs in the interior of a solid plate, where the two sides of the plate act as reflective surfaces. In a tunable filter, the interference arises in the air gap between the surfaces of two separate plates.

Tunable filter excitation map of NGC 1365 (Veilleux et al. 2003).

Fabry-Perot tunable filters offer several advantages over their single-plate counterparts.

  1. Tunable filters are tunable in wavelength space. By making small changes in the air gap between the two plates, one can change the transmitted wavelength.
  2. Tunable filters are tunable in bandpass. By making large changes in the air gap between the two plates, one can change the width of the transmission profile.
  3. Tunable filters provide very narrow bandpasses. Through rejection of significant sky and/or continuum emission (compared to conventional narrow-band filters), tunable filters can significantly improve observing efficiency for reaching a desired signal-to-noise ratio.
  4. Tunable filters can switch between two wavelengths on short timescales. Coupled with charge-shuffling on the CCD, this allows the user to correct for time-varying observational effects, including atmospheric transmission and sky brightness. It also allows for differential imaging on short timescales.

Alternatively, tunable filters function as wide-area, low-resolution spectrometers. From this perspective, a tunable filter achieves a significant efficiency advantage over conventional long-slit spectrographs (Jacquinot 1954).


Fabry-Perot tunable filters have been used for a wide range of applications, including (but not limited to) the following:

Fabry-Perot Physics

For general Fabry-Perot physics, refer to an optics textbook (e.g., Hecht's Optics) or this succint Wikipedia entry. For more details on tunable filters, see the references given below. We summarize here the important points.

Emission-line rings, viewed off the optical axis thru MMTF.

For Fabry-Perot etalons to operate as tunable filters, the plate spacing must be small. This ensures that the monochromatic spot at the center of the field of view is large, or equivalently the spectral resolution is low. The monochromatic, or Jacquinot, spot is defined to be the region over which the change in wavelength does not exceed √2 times the etalon bandpass. For a given wavelength, small plate spacing translates into low interference orders.

More specifically, the diameter of the Jacquinot spot is given by

  DJ = 2√(2√2) F / √(nN)
     ~ 3.3636 F / √(nN)

  where F = camera focal length (in pixels)
        n = order of interference
        N = etalon finesse (see discussion below)

For a given order, the monochromatic spot depends on wavelength only through the finesse. (The finesse, which we discuss below, depends on the transmission of the plate coatings.) However, as we discuss in the next paragraph, order is proportional to plate spacing d, so that DJ ~ 1/√d.

The transmitted wavelength at any point on the etalon is governed by interference principles: n λ = 2 d cos θ, where n is the interference order, λ is the wavelength, d is the effective plate spacing (including any optical effects due to the coating or air gap -- these may change with wavelength), and θ is the angle of the incident light with respect to an axis perpendicular to the plates. The transmitted wavelength can be modulated by either (a) changing the plate spacing d or (b) changing one's radius with respect to the optical axis, represented by cos θ. (At the CCD, the θ dependence translates into a radial dependence of wavelength; the optical axis is at the center of this radial pattern.)

Interference orders are separated by a free spectral range (FSR). Each transmission peak is characterized by an Airy profile. The ratio of the free spectral range to the instrumental resolution (in terms of the effective bandpass, defined as the profile area divided by its peak) is the effective finesse. In a perfect system, the effective finesse is equal to the reflective finesse, which is a function of the etalon plate coatings. However, irregularities in the plate surfaces and coating thicknesses, as well as deviations from perfectly parallel plates, degrade the effective finesse, and hence the efficiency of observing.

In summary, the interference equation governing the etalon is

  n λ = 2 d cos θ

  where n = order of interference
        d = plate spacing
	θ = angle of incident light with respect to the optical
        axis (θ = 0 at the optical axis)

The equations governing the instrumental profile are as follows:

The wavelength dependencies are summarized as follows:


Bland, J., & Tully, R. B. "The Hawaii Imaging Fabry-Perot Interferometer (HIFI)." 1989, AJ, 98, 723

Bland-Hawthorn, J., & Jones, D. H. "TTF: A Flexible Approach to Narrowband Imaging." 1998, PASA, 15, 44

Jacquinot, P. "The Luminosity of Spectrometers with Prisms, Gratings, or Fabry Perot Etalons." 1954, J Opt Sci Amer, 44, 761

Jones, D. H., Shopbell, P. L., & Bland-Hawthorn, J. "Detection and Measurement from Narrow-band Tunable Filter Scans." 2002, MNRAS, 329, 759

Taylor, K., & Atherton, P. D. "Seeing-Limited Radial Velocity Field Mapping of Extended Emission Line Sources Using a New Imaging Fabry-Perot System." 1980, MNRAS, 191, 675

Other Instruments with Tunable Filters

The TTF (Taurus Tunable Filter), formerly on the Anglo-Australian Telescope (see also articles listed above)

Fine Guidance Sensor - Tunable Filter, on the James Webb Space Telescope (see also these articles)

OSIRIS (Optical System for Imaging and Low-Intermediate-Resolution Integrated Spectroscopy), on the Gran Telescopio Canarias (see also these articles)

The Prime Focus Imaging Spectrograph, on the South African Large Telescope (see also these articles)

Instrument Specifications
System Components

The MMTF sits in the disperser wheel of the IMACS camera:

The optical path through IMACS and the MMTF.

Throughput curves for the old (black curve) and new, red-sensitive (blue curve) IMACS cameras. The new camera is more sensitive at all wavelengths, and 50% more sensitive at 7000Å.

The MMTF consists of 3 major components: the etalon and its support assembly, the CS-100 controller, and the cables connecting them.

A side view of the etalon, ensconced in its mechanical mount.

The etalon consists of two parallel plates composed of fused silica, each with a 150 mm clear aperture. The etalons are separated by a tunable air gap, with minimum and maximum separations of 1.3 microns and 10.7 microns. The plates are coated on the interior side (the side facing the other plate) with a reflective coating consisting of a multi-layer dielectric. (The finesse of the etalon is a function of the interior reflective coatings.) The exterior surfaces are anti-reflection coated. The plates are smooth at the λ level.

Capacitors and piezoelectric stacks lie near the plate edges. The voltages across the capacitors are used to measure the tilt and spacing of the plates. The stacks move the plates in three dimensions: a tilt around two axes (X and Y) and spacing along the Z axis. The tilt of the plates governs their parallelism; the best (most symmetric, and most flux in the peak) instrumental profile is achieved when the plates are most parallel. The spacing controls both the wavelengths transmitted and the width of the instrumental profile.

The CS-100 controller sends and receives electronic signals from the etalon to control the tilt and spacing of the plates. Once they are made parallel, the controller operates an electronic feedback loop to keep the etalon plates in position. The coarse values for tilt and spacing (Xcoarse, Ycoarse, and Zcoarse) are set using the dials on the CS-100 controller. The tilt and spacing can be changed in smaller increments using Xfine, Yfine, and Zfine values. The fine values can be changed using either the CS-100 dials or the IMACS control software. The observer should not change the etalon settings with the CS-100 controller dials.

Note: (X, Y)coarse and (X, Y)fine control the etalon parallelism. However, the coarse spacing Zcoarse controls the etalon bandpass, while the fine spacing Zfine governs the transmitted wavelength. The fine spacing can affect the bandpass, but only when changed by a large amount.

For more details on the CS-100, read the CS-100 documentation or a summary on the TTF website.

The etalon connects to the CS-100 through shielded cabling that runs from IMACS to the electronic cooling rack on one side of the Nasmyth platform.

The entire system was manufactured by IC Optical Systems Ltd.

Blocking Filters and Sensitivity

Observers have full freedom to tune the central wavelength of the MMTF bandpass within the range allowed by a particular blocking filter. The wavelength can be tuned from one exposure to the next, or within a given exposure (see the charge shuffling / frequency switching observing mode). A blocking filter is required to select only one transmission order (a given plate spacing transmits multiple wavelengths at a single location, each corresponding to a different order).

The currently available order blocking filters are listed below. Each has a bandpass intermediate between that of broad-band filters and that of the etalon. For each filter, we list the central wavelength and full-width at half-maximum in the IMACS beam, and provide a transmission curve. (The transmission curves were measured in a collimated beam, and are corrected to the IMACS f/2 beam; the correction is not exact.) We also provide a link to a table of etalon parameters at the relevant wavelengths, including FWHM values.

Intermediate Bandpass Order Blocking Filters
λ (Å) FWHM (Å) Transmission Etalon Parameters
5102 150 plot data table
5290 156 plot data table
6399 206 plot data table
~6600 ~260 plot data table
6815 216 plot data table
7045 228 plot data table
8149 133 plot data table
9163 318 plot data table

FWHM > free spectral range.

If you are considering the purchase of a narrow-band filter not listed here and are interested in making it available for use with MMTF, or have other questions about filters, please contact Sylvain Veilleux. Filter bandpasses must fall in the high-reflectivity part of the interior plate coating transmission curve; this range is roughly 5000-9500 Å (see Figure).

Transmission curve of the reflective etalon coating.

Changing filters between exposures is straightforward. However, each filter requires its own set of calibrations to determine the wavelength solution and etalon parallelism. It is thus important to plan ahead and find wavelength solutions for each filter during the day before a run. The observer is limited to the use of at most 2 filters in a single night because of calibration overheads.

Monochromatic Spot Size

For a given order, the area of approximately constant wavelength at image center, or monochromatic spot, depends on wavelength only through the finesse. For MMTF, the coating transmission is fairly constant above 5500 Å. However, the diameter of the spot does depend on plate spacing. We list here representative values for several different wavelengths and spacings:

Jacquinot Spot Diameters
Wavelength (Å) DJ (Zcoarse=-2) (arcmin) DJ (Zcoarse=3) (arcmin)
5100 13.81 10.71
6600 11.6 8.5
9150 11.5 9.5

1At 5100 Å the numbers are ~1'-2' higher due to a drop in finesse.


The available instrumental full-widths at half maximum are linked in the table above. For 6600 Å, the range is 6 - 13 Å. (The "effective bandpass", or the integral of the profile divided by its peak, is ~60% larger than the FWHM for a Lorentz distribution; the MMTF instrumental profile is approximately Lorentzian.) The FWHM values at other wavelengths, for a fixed plate spacing, scale with λ2. For a fixed wavelength, the FWHM is adjusted by making large changes in the plate spacing (d in the physical domain, Zcoarse in the electronic domain). In summary,

  FWHM(λ) ≅ λ2 / [2 d Finesse]

where the Finesse is mostly constant, except at the edges of the etalon's coating response. At the edges of the coating transmission curve (see Figure above), the reflectivity (and hence the etalon finesse) drops, broadening the profile and increasing its FWHM. Thus, at 5100 Å, the FWHM range is 8 - 14 Å. These values are slightly higher than at 6600 Å, while one would expect lower values for a constant finesse.

The bandpass can be changed during an observer's run only under exceptional circumstances.

Proposal Writing

Principal investigators are encouraged to contact Sylvain Veilleux with any questions/concerns that they may have about the feasibility of their project before submitting their proposal.

Strengths of the MMTF
Things to avoid with the MMTF
Proposal Checklist
Observing Strategies

NEW: In early 2009 a new method of observing with the MMTF was tested and found to drastically reduce the uncertainty in plate spacing and thus cut nightly overhead costs by up to 1-2 hours! By forcing IMACS (and thus the MMTF) to remain at a roughly constant gravity angle, the etalon will remain parallelized throughout the entire night, eliminating the need for costly re-parallelization and reducing the drift in the wavelength-spacing solution. Be sure to follow the directions below in order to take full advantage of this revolutionary technique.

Computer Controls

MMTF observing is done using the standard IMACS data acquisition software along with a set of special built-in procedures. These include the ability to set the etalon fine X, Y, and Z values in the Hardhat GUI, as well as a set of programmable scripts that are run from the MMTF version of CamGUI. These scripts coordinate commands to the camera and the CS-100 for the following observing and calibration tasks:

Note that all access to etalon X, Y, and Z settings is through the IMACS software; the observer should not change the etalon settings with the CS-100 controller dials.

Before the Run
  • Choice of observing mode.

    The MMTF is intended to be operated in one of three modes:

    1. Staring mode. In staring mode, the etalon gap is fixed and the result is a narrow-band image in a single wavelength at each point in the field. However, the wavelength varies slowly across the field of view.
    2. Wavelength scanning mode. In scanning mode, an image is taken at a series of different etalon spacings. The result is a low-resolution spectrum at each position.
    3. Charge shuffling + frequency switching mode. See this section for a description of the charge shuffling and frequency switching mode. If you intend on using the charge shuffling mode, you should contact Sylvain Veilleux for planning assistance.
  • Choice of Zcoarse. The Zcoarse, or coarse plate spacing, choice is a balance between the size of the monochromatic spot and the signal-to-sky contrast. Larger values of Zcoarse yield somewhat smaller monochromatic spot sizes. However, they also have smaller bandpasses, which means that the amount of sky emission transmitted is lower, thus increasing the signal-to-noise ratio.
  • Choice of wavelength and filters. The available filters are listed in this table. The observer is limited to the use of at most 2 filters in a single night because of calibration overheads.
During the Afternoon

For each filter to be used during the first night, the following steps should be taken. For each step, refer to the Calibrations section for more details. Be sure that these calibrations are all performed at a gravity angle of Θgrav ~ 0° (ΘNASW = -47.3°).

  1. Parallelize the etalon. [procedure]
  2. Capture a reference image of an emission-line ring . [procedure]
  3. Take a data sausage or a sausage link. [procedure]
  4. Prepare your target list. [sample]
At the Start of the Night

For each step, refer to the Calibrations section for more details.

  1. Acquire the first target, and ensure that the field of view is properly rotated (Θgrav ~ 0°).
  2. Capture a reference image of an emission-line ring, and compute and apply a correction to A and Zfine [procedure].
  3. Compute and set the etalon to the appropriate Zfine for the wavelength of interest using WCALC.
Throughout the Night
  1. Before each exposure, direct the telescope operator to re-acquire your target coordinates, in order to rotate IMACS to Θgrav ~ 0° (be sure that the rotation mode in your target file is "GRV" and the offset is "-47"; see the following example). When doing a sequence of multiple exposures, re-acquire before each one in order to maintain a roughly constant gravity angle throughout the night. This will minimize errors in your wavelength solution and drastically cut down on nightly overheads by up to 1-2 hours.
  2. Take a reference ring regularly to check for drift in A (wavelength zero-point) and the line profile (parallelism) (procedure). This should be done every ~30 minutes while the temperature is changing dramatically at the start of the night, or every hour when the temperature is stable. If the line profile has degraded in quality (less narrow and/or less symmetric), you may have to recompute the parallelism (procedure). Use this reference ring to apply a correction to A and Zfine, using WCALC, to achieve the wavelength of interest.
  3. Exposure times should be no longer than 20-30 minutes to allow frequent checks for wavelength zero-point drift. Drift occurs more rapidly at the beginning of the night, due to more rapid temperature changes at this time. If you are doing on- and off-band imaging, be sure to do the on-band immediately after checking for wavelength drift while the wavelength solution is valid. The exact wavelength of the off-band image is less important and, thus, should always be done second.
  4. For exposures where it is important to achieve a continuous image uninterrupted by features or chip gaps, dither by 30 arcseconds or more in each direction. Dithering is important for IMACS imaging in order to bridge the inter-chip gaps, which are ~50 pixels, or 10 arcseconds, on average. The MMTF also introduces a set of reflections of the chip edges (of relative intensity a few percent) which are not corrected by the flatfields. These reflections contaminate ~100 rows and/or columns of each chip. Dithering over 5' scales is required for observations of extended sources or fields in charge-shuffle mode.
At the End of the Night
  1. Get dome flats at each wavelength that you observed during the night [procedure].
    (Can also be done in the afternoon prior to observing)
  2. Get a series of biases (0-second exposures).
Data Reduction

The following sections describe the MMTF pipeline for reducing IMACS data. This collection of perl-driven IRAF scripts is a work in progress and still have not been tested exhaustively - they should be used with care. Please contact Michael McDonald with comments, questions, or bug reports.

Download the script package v1.4 (post April '08 data), or package v1.1 (pre April '08 data)

Download the appropriate optical axis file


Additional tasks:

The data from IMACS emerge as 8 individual files for each exposure, one file per chip. The bias subtraction, flatfielding, cosmic ray masking, sky subtraction, image registration and combiningg steps described below operate on these individual files.

Bias Subtraction and Flatfielding

The bias subtraction and flat fielding steps are straightforward. To use the script, make a list of bias files, one per line, disregarding the chip and ".fits" labels. (Alternatively, copy the bias files from a previous reduction into the current directory and select the "--bias" switch.) In a second file, make a list of each source file, one per line. On each line, the source file is followed by a comma-separated list of flats to be combined and applied to that exposure. Alternatively, one may leave a blank, in which the flats default to those from the previous exposure.

For example:

Contents of file "bias.lst":


Contents of file "objfl.lst":

ccd0011 ccd0090,ccd0091,ccd0092
ccd0013 ccd0093,ccd0094,ccd0095

In this case, frame 0011 and 0012 share the same flats, as do ccd0013 and ccd0014. The script is then run:

% imacsbf bias.lst objfl.lst

This produces an IRAF script, which can then be run from the IRAF prompt. The following steps are performed by the IRAF script:

  1. The row overscan is subtracted from all files.
  2. The biases are combined. The combined bias has filenames biasc1.fits, biasc2.fits, etc.
  3. The combined bias is subtracted from each source and flatfield file.
  4. The column overscan is subtracted.
  5. The flats are combined. The combined flats have filenames flat1c1.fits, flat1c2.fits, . . . , flat2c1.fits, etc.
  6. Each flat is normalized so that the central region of the FOV has an average value ~1. The same normalization factor is applied to each chip.
  7. If specified, scattered light signatures are removed from the flat field.
  8. The source files are flatfielded.
Cosmic Ray and Bad Pixel Masking

Cosmic ray removal is achieved in two stages. First, we attempt to remove cosmic rays for each image individually. This is done using Pieter van Dokkum's LA Cosmic task for IRAF (for details, see: LA-Cosmic). This task is called in the imacscr script, which has switches to allow control of such things as the sky level and the threshold sigma for rejection. Generally, four iterations is enough to remove the majority of the cosmic rays in each field. A typical run of imacscr is as follows:

Contents of file "mask.lst":


Then run the following command:

 % imacscr mask.lst 

This will produce an IRAF script ( which should be run from the IRAF prompt. The following steps will be performed by

  1. The LA Cosmic task will be loaded into IRAF using the stsdas.task command and the default values such as readnoise and gain for the IMACS camera will be set.
  2. LA Cosmic is run on each of the 8 chips for each image prefix.
  3. Known bad pixel regions (e.g. bad columns) are appended to the cosmic ray masks, which are saved as msk*.fits

It is important to note a few things about this routine. First, LA Cosmic runs best on images that have not had the sky subtracted. We therefore recommend running imacscr immediately after imacsbf. However, if you have already subtracted the sky for some reason, you can specify roughly what the sky was using the --sky=[value] switch. Secondly, imacscr will produce output images with cosmic rays removed, saved as cr*.fits. These images are useful to check how well LA Cosmic is working, but will not be used by our pipeline (more on that later!). Finally, although a very effective task, LA Cosmic tends to leave residual bright pixels unmasked (although very rarely does it miss a cosmic ray completely!). This is remedied by using the IRAF "crgrow" task later on to expand the masks by a single pixel in each direction.

Sky Subtraction

Sky subtraction is performed using the MMTF calibration software. Given knowledge of the mapping of the CCD pixels to the focal plane and the location of the optical axis, the sky spectrum is azimuthally averaged. Sources and cosmic rays are filtered out using a biweight statistic. If necessary, median smoothing is employed to smooth out the sky spectrum. This results in a robust removal of the sky using a single command.

To subtract sky from one or more exposures, create a file containing, on each line, the exposure name, the (approximate) filter wavelength, and the optical axis coordinates (chip #, column #, row #). For instance:

Contents of file "skysub.lst":

fccd0011  6600  5 -12 4056
fccd0012  6600  5 -12 4056
fccd0013  8150  5 -12 4056
fccd0014  8150  5 -12 4056

Then run the following command:

% ringbatch --list=skysub.lst --sub --ms=1

For single exposures, one may simply run the RING command:

% ring ccd0011 6600 --sub --ms 1

This procedure runs non-interactively. On a multi-threaded 3 GHz PC processor running Linux, with 1.5 Gb of memory, the sky subtraction takes approximately 30-45 minutes per exposure (for unbinned data). The result is a series of files with the "ss" suffix prepended. The default FOV has a radius of 4000 pixels in the focal plane; pixels outside this radius do not have the sky subtracted from them. These pixels can be masked using the procedure described below.

If the optical axis is not accurate to within ~20 pixels or so, dipole residuals may appear after sky subtraction, either in individual exposures or in a stack of several exposures with small dithers. If the etalon is properly parallelized, the optical axis estimate provided by the instrument scientist should be correct at this level. In the event that it is not, we recommend that the observer re-measure the optical axis using the actual science data and bright stars in the field of view, following this procedure (except that the mask image/primary ghost pairs are in this case stellar image/ghost pairs). Note that the stellar ghosts will be out-of-focus.

For exposures with large, extended sources in the field of view, these sources may have to be masked before accurate subtraction can be performed. Using the --skip flag, the user can ignore specific chips in the azimuthal calculation. If an extended source is at the center of the field of view, the observer must extrapolate the sky spectrum (as output by the RING/RINGBATCH procedures) into the central regions. The user may modify the spectrum by changing the biweight values in the *.rings output file, using the good part of the spectrum as a guide. This modified file (which must have the same format as the original) can then be put back into RING using the "--spec" command-line switch. (This feature is not currently available under RINGBATCH.)


Median astrometry errors for the new IMACS camera. The upper panel shows absolute errors (in arcseconds) as a function of pixel position, while the lower panel shows the magnitude (multiplied by 500) and direction of the errors as a function of pixel position. These fields were constructed using 8000+ stars over several different pointings.

Perhaps the most challenging aspect of IMACS data reduction is determining the astrometric solution that maps the physical CCD coordinates to locations on the sky. Because of the small f-ratio of the IMACS camera (f/2) when the MMTF is in use, as well as the wide field of view, significant deviations from a linear mapping occur. If your science data occupies only a small, central region of the FOV, standard astrometric techniques may be applied on a chip-by-chip basis. However, to rectify the whole field, we suggest the following procedure.

The higher-order polynomial terms of the FOV-to-sky mapping are known from observations of dense star fields. We provide a script that applies this mapping to any IMACS data, using estimates of the tangent point and sky position angle from the FITS headers. This script, imacswcs, is executed in the following way:

Contents of file "wcs.lst":


Then run the following command:

 % imacswcs --list=wcs.lst 

For single exposures, one may simply use the following syntax:

 % imacswcs ssfccd0011 

This will produce an IRAF script ( which should be run from the IRAF prompt. When executed, this script will apply the WCS information to all 8 chips for each exposure. This information contains the high-order distortions in the chips as well as the accurate chip spacings and the rotation and scale of the field of view. This information, which is only stored in the header for the time being, will be applied in the next step by the imacsreg script.

Registration and Image Mosaicing

In order to properly register each field of view, imacsreg must be run. This script uses the high-order solution provided by imacswcs and applies various IRAF tasks to fine-tune this solution on a case-by-case basis. The script is run as follows:

Contents of file "reg.lst":


Then run the following command:

 % imacsreg reg.lst 

This will produce an IRAF script ( which should be run from the IRAF prompt. The following steps will be performed by

  1. Cosmic ray masks are extended by one pixel in each direction to prevent ringing caused by residual cosmic rays in the interpolation steps.
  2. Masked pixels in the input files (e.g. ssfccd0011) are replaced by 0.0 (the sky value) to prevent blurring of cosmic rays. These values will not play a role in the final image, since we will be ignoring contribution from masked images. The only reason for this coarse interpolation is to prevent ringing in the next step.
  3. The IRAF task "mscred.mscimage" is run on each individual chip in order to apply the high-order TNX corrections. The bad pixel masks undergoes the same transformation as the image. This task is run on each chip individually because "mscimage" will not apply TNX corrections on mosaiced images (this is a known IRAF issue).
  4. The chips are combined into a composite MEF file and "mscimage" is run a second time to create a single file with a single bad pixel mask including cosmic ray, bad column and chip gap locations.
  5. A catalogue of stellar positions is created using the IRAF task "mscred.mscgetcat". Stars falling anywhere within roughly 3500 pixels from the center and with magnitudes ranging from 19-10 are used.
  6. The IRAF task "mscred.msccmatch" is used to match the stellar positions to pixel positions in order to determine the lower-order corrections for the mosaicced image. These corrections will be different from exposure to exposure due to various atmospheric effects and pointing errors. These small corrections are also applied to the bad pixel masks.
  7. Using the catalogue of stellar positions, a catalogue of PSFs are made for each field of view. This will be used in the combining stage when PSF matching is required.
PSF-Matching and Stacking

Before combining dithered exposures into a single, deep image imacscombine will PSF-match and align the image groups. The script operates on several groups of images, each group occupying a single line and separated by commas. imacscombine, which is executed once imacsreg has produced mosaiced images with accurate astrometry, is run as follows:

Contents of file "combine.lst":


Then run the following command:

 % imacscombine combine.lst 

This will produce an IRAF script ( which should be run from the IRAF prompt. The following steps will be performed by

  1. Using the PSF catalogues produced by imacsreg, the mode seeing value is computed for each exposure and inserted into the header.
  2. The images are aligned based on the WCS information. These offsets account for any dithering or pointing errors. Bad pixel masks are also shifted accordingly.
  3. All similar frames are degraded to the worst seeing using the iraf task "psfmatch". **NOTE** If one image has anomolously bad seeing, it should be removed from the list to avoid degrading the images too far.
  4. Similar frames are stacked using the "imcombine" task in IRAF. The process of combining images will ignore masked pixels and will also reject all values above a certain threshold (default=2.0 sigma), thus removing any residual cosmic rays.
Converting to Multi-Extension FITS (MEF) Files

The routine provided creates an IRAF script that will convert the eight individual files in an exposure, one per chip, into a single multi-extension FITS file. Such files are useful for input to, e.g., the IRAF mosaic reduction (MSCRED) package. MSCRED can, among other things, display the entire mosaic using the task MSCDISPLAY, or refine the image WCS for astrometric calibration.

Radial Masks

Radial masks can be created using the task RADMASK, part of the MMTF calibration software. This procedure assumes knowledge of the optical axis. The masks created can have any inner or outer radii, and could be used for cosmetic masking of the outer, unilluminated parts of the field of view or to mask an extended object centered in the FOV.

Optical Axis File (optaxis.dat)
The position of the optical axis can vary slightly with each mounting of the MMTF in the IMACS disperser wheel. This table contains measurements of the optical axis position which are required to properly subtract the sky. For a description of how to measure the optical axis position, see this page.
Date Created Optical Axis File
Mar 2010 optaxis_0310.dat
Nov 2009 optaxis_1109.dat
Jul 2009 optaxis_0709.dat
May 2008 optaxis_0508.dat
Dec 2008 optaxis_1208.dat
Nov 2007 optaxis_1107.dat