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OSU College of Arts and Sciences Department of Astronomy

The TIFKAM User's Manual

Contents

Introduction
System Overview
System Characteristics and Performance
The HAWAII-1R HgCdTe Detector Array
Observing Setup
Observing Techniques
Calibration
Additional Information

Introduction

TIFKAM is a multi-purpose infrared imager and spectrometer designed for operation at 0.9-2.5 micron wavelengths. The instrument currently uses a 1024x1024 HAWAII-1R HgCdTe detector array built by Rockwell and purchsed for TIFKAM by Dartmouth using funds provided by an NSF MRI grant. Selectable re-imaging cameras provide a choice of two plate scales and a variety of spectroscopic resolutions. Internal mechanisms control the selection of filters, dispersers, focal-plane masks, and cameras. All mechanisms are user controlled and can be used to rapidly reconfigure the instrument, which allows for substantial observing flexibility.

Why "TIFKAM"?

When originally designed, the instrument was named "MOSAIC", an acronym for "MDM/Ohio State Active Infrared Camera". The intent was to have an internal tip-tilt system to correct out first-order image motion ("active optics") to improve image quality. The tip-tilt system was never implemented. We were then loaned a 512x1024 ALADDIN InSb array by Kitt Peak National Observatory, in return for letting the KPNO community use the instrument at the KPNO 4m and 2.1m telescopes. The problem was that MOSAIC was also the name of the 8Kx8K CCD imager at KPNO, which caused confusion. Mike Merrill at KPNO suggested the name "The Instrument Formerly Known As MOSAIC", hence TIFKAM.

Aren't you glad you asked?

System Overview

TIFKAM is a general-purpose 0.9-2.5 micron imager and moderate-resolution spectrometer built as collaboration between Ohio State and the MDM Observatory. The instrument uses a 1024x1024 HAWAII-1R HgCdTe detector array. Internal cameras provide three choices of pixel scales (see below) or a pupil-viewing camera.

There are two 9-position filter/grism wheels in TIFKAM providing a variety of filters and spectroscopic modes. There are 4 grisms available, in approximate order of increasing dispersion:

The population of the instrument filter, slit, and camera wheels changes occasionally, but generally includes standard broad band imaging filters (i.e. JHK) and a selection of grisms and slits.

The Light Path Through the Instrument
 

The control interface for the instrument is the Prospero package written at OSU. More information on Prospero can be found on the Prospero web page.

Please note that Prospero is not an image processing package, but an interactive instrument control system. Simple commands for inspecting images are included in the Prospero command suite, but the observer is expected to use their favorite package to fully reduce and analyze their data (e.g. VISTA, IRAF, IDL, FIGARO, MIDAS, ZODIAC, HP-65, ABACUS, FINGERS). A useful summary of the Prospero commands pertinent to TIFKAM is provided in the form of a 2-page quick-reference card.

The instrument has many operational modes and has not been extensively tested in each. There are also several known problems with the instrument which we note below. Our intent is to work on them as soon as possible, usually the next available opportunity to warm up and open the dewar. For historical interest, we also provide a list of problems the instrument has had in the past.

System Characteristics and Performance

Imaging

TIFKAM was designed to deliver good image quality over the entire field of view of the 1024x1024 detector array with reasonably high throughput.

TIFKAM has three camera lens sets (f/5, f/7.5, and f/16.6) with the following pixel scales:

Camera 2.4m Hiltner 1.3m McGraw-Hill
f/5 0.30 "/pix0.55 "/pix
f/7.50.20 "/pix0.37 "/pix
f/16.60.05 "/pix0.09 "/pix

A fourth camera is used to view the pupil and rarely used by observers. The f/16.6 camera is primarily intended for spectroscopic use or to take advantage of periods of unusually good seeing at the 2.4m. Most observers taking direct images will use either the f/7.5 or f/5 cameras.

Spectroscopy

Long-slit spectroscopy is enabled by grisms inserted into the beam by the filter wheel. Below is a table that summarizes the resolutions and wavelength coverages of the various possiblities.
GrismBlocking
Filter		 Resolution
(2 pixels)		 Wavelength
Range (nm)		 Illuminated
Rows
J/KJ14501110 - 1360 (2nd order)400 - 982
J/KK13251970 - 2420  (1st order)160 - 704
HH11501480 - 1800 (2nd order)470 - 918
J+HIJH720950 - 1910 (1st  order)1 - 900
JHKJHK7501220 - 2490 (usable to 2290; limited by blue leak)1 - 1024
Note that the actual wavelength coverage will depend on the filter selected in the other filter wheel (i.e. if using the J/K grism and the K filter as a blocker, the wavelength coverage will only be 2000 to 2400nm due to the transmission of the filter); we are currently trying to obtain special filters to use specifically as spectroscopic blockers. The table above gives the resolution assuming the 54-micron slit is also inserted into the beam; this corresponds to roughly two pixels FWHM for unresolved lines. There is also a 100 micron slit that will give four pixels per resolution element and correspondingly lower spectral resolution. Note that this slit is recommended when the seeing is poor.

The HAWAII-1R HgCdTe Detector Array

TIFKAM Detector Characteristics
Full Format1034x1024
Pixel Size18 microns
Active Region[11:1024,6:1019] (1014x1014)
System Gain2 electrons/ADU
Read Noise15 electrons RMS
Full-Well Capacity~105 electrons
Minimum Integration Time4.29 s
The HAWAII-1R array has a series of reference pixels appear as a dark strip around the outside margins of the raw image. While these look superficially like CCD overscan pixels, they are instead physically separate pixels in the CMOS readout multiplexer that are not connected to the HgCdTe detector layer, and insensitive to IR light. The actual active area of the detector is 1014x1014 pixels with 5 lines of reference pixels above and below the active area veritically and 10 columns of reference pixels on either side.

Dark Current and Read Noise

The dark current in the array is low. We currently measure <0.4 ADU/pix/sec dark current, some of which may be due to background radiation in the dewar proper. The dark current will be higher if the instrument has been cold for less than 36 hours.

The read noise of the array is ~15 electrons.

Maintenance

The only maintenance task required during an observing run is to keep the LN2 topped off. Our experience is that filling the instrument once per day is sufficient.

Linearity and Saturation

The array becomes significantly non-linear before the full-well capacity of the detector is reached. Further, an accurate non-linearity correction depends on the signal rate as well as the total signal collected. However, non-linearity is <1% for <11,000 ADU at signal rates ranging up to ~600 ADU/sec (about the background rate at K on a warm night with the f/7.5 camera). We are working on obtaining more information about the reproducibility and signal-rate-dependence of the non-linearity; we currently suggest that interesting signals be kept to <8000 ADU so that non-linearity corrections should be <0.2%.

There may be a small offset between the integration time requested by the observer (and entered in the image header) and the actual time interval between the two reads of the double-correlated sampling cycle. This will result in an apparent non-linearity in the measurement of a constant signal as a function of integration time. Until we understand this better, we advise against the use of bright standard star observations at very short integration times to photometrically calibrate object fields at much longer integration times.

Hard saturation of the detector array occurs at about 25,000 ADU. The array becomes seriously non-linear (>5%) at about 17,000 ADU.

Bad Pixels

There are many dead, hot, unresponsive, and just plain bad pixels on the array.

After Image or Residual Charge

The detector array in TIFKAM exhibits a residual signal from bright sources. The magnitude of the residual image seems to depend on the brightness of and total signal recorded from the source. For typical observations the residual will be 0.5-2% of the originally detected signal. Reading the array several times reduces the magnitude of the residual image to <<1% of the original signal.

We suggest that if residual images will seriously compromise your results, you dither the telescope faithfully and, perhaps, read the array several times (set the exposure time to 0 seconds, execute an mgo 3 command, and throw those images away) between each science exposure.

Observing Setup

Focusing

Both the internal optics and the telescope should be in focus for the most optimal images.

The internal optics can be focussed by inserting one of the spectroscopic slits into the beam and then running the camera focus through a range of values. Approximate values (June 2008) are 

    f/7.5  - camfocus 1340
    f/5    - camfocus 2730
    f/16.6 - camfocus 400
The focus curve for the f/16.6 camera is very shallow, so if remeasuring you should use big steps (200) in camfocus, whereas for the faster cameras steps of 50-100 are best. For spectroscopy, we recommend that the optimum camera focus be determined using a calibration line source. The value of the the various camera foci do not change unless the detector array is removed or the optics are realigned, both of which require opening the dewar and are (hopefully) rarely done.

The telescope can then be focused in the normal manner. We generally find a rough focus by starting movie (which reads out the array continuously) and rapidly running the telescope focus in and out until the image looks good, then taking a series of images at different focus settings.

Note that all the filters and prefilters are in near-collimated light, so there is no change of optimal focus with wavelength.

Observing Techniques

Imaging Observations

There are five (5) mechanisms that you control from within the Prospero program: a slit/focal plane mask wheel ("slit"), two filter wheels (called "filter" and "prefilt"), the choice of camera ("camera"; which sets the imaging plate scale and spectral resolution), and the internal focus position of the camera ("camfocus"). The slit, filter, and camera selections can be most conveniently found by typing print slit or print prefilt etc., to list the populations of that particular mechanism.

The command syntax for moving these mechanisms is 

   [mechanism] [value]
For example, to set up for J band imaging with the PK-50 blocker, for example, one would execute 
   prefilt 2
   filter 6
in any order.

The slits are roughly 480 pixels long. The 54-micron pixel slit is slightly rotated with respect to the array, by about 2 pixels over the length of the slit, whereas the 100-micron slit appears aligned with respect to the array to within 1 pixel.

There are three ways of taking exposures:

go takes a single exposure
mgo n takes n integrations and writes n frames to disk
avego n takes n integrations and writes 1 averaged frame to disk
The array becomes significantly non-linear beyond ~8000 ADU that it is advisable to keep the integration times short enough to keep the background less than this level. At this point the background noise overwhelms the read noise anyway. The upper right quadrant seems to become non-linear at a somewhat lower level than the rest of the array.

Spectroscopic Observations

Use imaging mode to acquire spectroscopic targets, offset the telescope to move the target onto the array location of the slit (which can be found by taking an image of the slit against the night sky; without the grism in place of course), acquire a guide star to "lock" the telescope onto the object, put in the slit, check that the telescope did not move while acquiring a guide star by taking another image (and correct the position if it did move), put in the grism, and start taking data. Note that the first few spectra may have residual image artifacts due to the previous imaging observations. Executing several short "dummy" observations prior to taking science data will greatly reduce the magnitude of these artifacts. We recommend taking spectra of point sources at several positions along the slit to provide sky measurements and to aid in removal of systematic effects such as bad pixels which may occur in any one spectrum. Because the slit is so narrow, we strongly recommend the use of the autoguider at both the 2.4-m and 1.3-m telescopes.

NOTE: Although the grisms are in the nominally collimated beam, we have observed that the lines from a comparison source are in best focus at a camfocus value of approximately 250, whereas a direct image of the slit is best focused at a value of approximately 100. We recommend that the optimum camera focus for spectroscopy be empirically determined using a comparison lamp spectrum.

Telescope Focus

There is a fairly strong focus dependence of the telescope on the ambient air temperature at both the 2.4-m and 1.3-m telescopes, and the 1.3m telescope can go out of focus at large zenith distances.

Prospero Scripts

Although Prospero commands are normally entered from the terminal, it is possible to execute a list of commands stored in an external text file. These scripts make it possible to execute complex or repetitive observing sequences, including loops and conditional branching from the vocabulary of individual Prospero commands. These may be written and stored on disk prior to observing. The Prospero Command Procedure Scripts Manual provides comprehensive coverage of this subject.

Calibration

Standard Stars

Photometric Standards

The best near-infrared standards are those defined by Elias et al (1982, AJ, 87, 1029), but these stars are all too bright to observe with TIFKAM without neutral density filters (i.e. they saturate the detector array in the minimum available exposure time). There are several sets of fainter standards including those measured by Carter & Meadows, the UKIRT Faint Standards, and a set of stars being measured to support NICMOS. The Carter & Meadows measurements appear to be excellent quality, but the stars are relatively bright (e.g. K = 9-10 mag) and may not be observable with TIFKAM without a neutral density (ND) filter. Since we do not know how "neutral" an ND filter is in the near infrared, this makes this set of standards somewhat problematic. Note, however, that an ND filter is currently installed in TIFKAM . The UKIRT standards are probably fine for measurements requiring no better than 5% accuracy or so.

We suggest that the list of standards prepared for NICMOS (Persson et al. 1998, A.J., 116, 2475) be used when the most accurate photometry is desired.

Spectroscopic Standards

The near-infrared region has many strong absorption features due to various molecules in the atmosphere. One of the best ways to remove these features is to ratio the spectrum of the target object with the spectrum of a featureless source observed with the same instrumental setup and airmass. In the J and K band, A stars provide good atmospheric standards, since they have only H absorption features at 2.17 microns (Br-gamma) and 1.28 microns (Pa-beta). In the H band, A stars have many Br-series absorption features that make them somewhat problematic. Note that Kurucz models of A stars actually seem to reproduce the near-infrared spectrum of these stars reasonably well, so these may be used to correct for the intrinsic absorption in the stellar atmosphere.

Wavelength Calibration

In the JHK region, the OH airglow lines which are such a nuisance do have the virtue of providing a useful grid for spectral calibration. However, at both the 2.4-m and 1.3-m telescopes, the spectral comparison sources built into the MIS boxes provide an excellent Ar spectrum, with a few He and Ne lines thrown in. One may obtain a high S/N calibration spectrum with 30 to 60 sec exposures (followed by an equal dark observation for bias subtraction) using the lamp sources, so this technique is recommended.

Flat Fielding

The HgCdTe detector array is reasonably flat; we have measured ~1% photometric accuracy for standard stars (at K) without flat-fielding the data. The best flat fields seem to be made by observing the dome flat field screen with the lights on (at a low level in imaging mode) and with the lights off, forming the flat from the difference of the two. This technique seems to remove any significant scattered light or thermal emission component from the resulting flat.

Additional Information

Web Pages:

TIFKAM Web Page:

http://www.astronomy.ohio-state.edu/MDM/TIFKAM/
This site includes descriptions of the instrument (with pictures), and access to copies of all of the documentation.

Prospero Support Web Page:

http://www.astronomy.ohio-state.edu/~prospero

This web site includes a searchable online manual for the Prospero package, and copies of all of the documentation. A help line and other services will be included in the future.