The IRIM User's Manual

The KPNO Infrared Imager (IRIM) is a general-purpose instrument for moderately wide-field imaging in the 1 - 2.5 micron region using a 256 X 256 HgCdTe NICMOS3 array. It is used primarily for broadband work in the J, H, K and K' filters, but it can also be used with narrowband filters for special applications.

IRIM is an uplooking cryostat which interfaces to the f/15 foci of the 2.1-m and 4-m telescopes. The telescope focal plane is warm, about 5 cm above the entrance window. The dewar carries analog and digital electronics boxes and communicates through optical fibers with the instrument computer. Apart from an "on" button for the dewar electronics, there are no external user adjustments. Cooling is by LN₂ with a dual reservoir system. The large outer tank cools the radiation shields, optics, and other innards, while the inner tank cools the detector mount assembly. Cryogen refill can be done with the instrument in place on the telescope. Hold time is typically 12 hours, but a refill halfway through a long winter night is advised for caution's sake.

Optical Diagram

An optical diagram of IRIM is shown below.

The internal optics are a refractive, collimator-camera design with a cold pupil stop of fixed size. The telescope focal plane (FP) lies approximately 75 mm above the CaF₂ window (W). The collimator (L1) images the telescope pupil onto the cold stop (CS). The 8-position filter wheel (FW) is located very close to the cold stop and carries 7 filters and a cold dark slide. The camera (L2, L3, L4) then reimages the focal plane onto the array (D). The overall optical reduction of 4:1 was designed to give field of view, particularly at the 4-m telescope, at the expense of optimal spatial sampling of the image.

Field and Filters

Measurements at the 4-m show a wavelength dependence of scale: 0.608 arcsec/pix at J, 0.605 at H, 0.603 at K and narrowband filters. Proportional changes should occur at the 2.1m, but they have not been determined accurately.

IRIM is operated by the user from the SUN workstation in the telescope control room, through the WILDFIRE system, a transputer based system which communicates over optical fibers. WILDFIRE supports fast co-adding in place, movie mode, and data transfer directly to the SUN.

The WILDFIRE system uses transputers and transputer links to control and acquire data from IRIM. A transputer is a single-chip microcomputer with its own local memory and communication links, which can operate either by itself or in conjunction with other elements linked to form computing arrays and networks. The WILDFIRE system consists of three main hardware components:

• The IRIM instrument control unit inside the DCU (Digital Control Unit) box contains two transputers which provide housekeeping data and control and generate the sequences which operate each array. The motor controller module is also inside the DCU, eliminating the need for a separate motor controller box. All motor, data, and power cables are connected to the DCU.
• The DSP unit (a VME based Digital Signal Processor system located inside the black Heurikon box in the computer room) contains eight transputers which provide the math processing needed to do coadding as the data are taken and buffer space for finished data before it is transferred to the SUN computer.
• The B016 unit (a programmable dual-port memory and interface board located inside the black Heurikon box in the computer room) interfaces the transputers and the SUN and handles the formatting of data before it is saved to disk.

Communications between IRIM and the DSP take place over transputer links implemented on an optical fiber cable. The B016 interconnects the transputer DSP to the SUN SparcStation computer via a VME to SBUS converter within the Heurikon box.

The WILDFIRE user interface on the SUN is implemented within the TCL (tool command language) environment. The data appear as IRAF images, produced (in IEEE 32-bit floating point format) via IMFORT routines so that they can be manipulated and archived to tape inside IRAF. It is important to note that these images are NOT PROTECTED in any way and can be overwritten if the full path names of existing and new images are the same. The data may be written to Exabyte or DAT on local tape drives or sent via 'ftp' to one's home institution. Depending on the amount of header information, a single FITS file of a 256 X 256 image is about 270KB.

At each of the telescopes where IRIM is used are two SUN workstations for data acquisition and reduction. Under the present version of WILDFIRE, the workstations khaki (4-m) and royal (2.1-m) are used for data acquisition. The other workstations [pecan and lapis at the two telescopes] have common access to the data disk, so additional observers can reduce and analyze the data independently. A third SUN serves as the telescope control, with a terminal at the LTO station; a hardwire link between the TCS and instrument control computers is used to send TCS commands to the telescope (singly, or within TCL scripts) and to retrieve telescope information for the image header. A schematic depiction of this arrangement is shown in Fig. 2.

Important Note: The disks within the primary workstations khaki and lapis are designated /data1. At the 2.1-m, WILDFIRE is run on the secondary workstation royal, whose partition is /data2. The disks are cross-mounted so that access to both is possible from either machine. However, such cross-accessing (e.g., /data1 from royal) is significantly slower than accessing the disk resident in the workstation. Therefore, it is imperative that the partition used for storing data taken by WILDFIRE be /data2 on the 2.1-m telescope! While it is possible to designate /data1 as the WILDFIRE data partition, operation will be much slower and subject to crashes, so don't do it. At the 4-m, one may designate either /data1 or /data2 as the data partition.

The array has a rather nonflat response, with \(+- 25% variations in response and about three full high-low cycles across the field.

The number of "bad" pixels is somewhat a matter of definition:

• There are about 250 nonresponsive (open or shorted) pixels scattered uniformly across the array, typically singly or in 2 X 2 clumps.
• There are about 250 pixels with noisy dark current; i.e., a standard deviation calculated from a series of darks is > 2 sigma above the mean for the whole array.
• There are about 900 pixels with response outside the range mean +- 25% or equivalently mean +- 2 sigma.
• There are about 1000 pixels with noisy response; i.e., a standard deviation calculated from a series of flatfield exposures which is > 2 sigma above the mean for the whole array.

Clearly some pixels fall in more than one category, and a useful definition of "bad pixels" will depend on the program. For most IRIM programs, the combination of categories (3) and (4) is useful. A bad pixel map can be constructed interactively using IRAF tools and applied during the reductions. A table of the detector characteristics appears on the following page.

Linearity and Dark Current

Both InSb and HgCdTe detectors utilize a hybrid architecture in which each pixel has an associated unit cell which controls the biasing and readout of that pixel. Thus, each pixel is essentially independent of the others, and effects seen in CCDs, such as charge bleeding or trailing from saturated pixels, are not present. However, this independence also means that such properties as linearity and dark current can vary from pixel to pixel, and it is necessary to calibrate these effects for optimum scientific performance. When a pixel is reset, the voltage difference (bias) between the pixel and detector substrate creates a depletion region which acts as a potential well for the collection of (mostly) photogenerated carriers. Electrically, one may consider this potential well as a capacitor. As charge accumulates in the pixel, the depletion region fills in, increasing its capacitance and that of the entire pixel node. Coupled with the steadily decreasing bias on the pixel, this yields a sublinear voltage-charge relationship, which quickly rolls off (saturates) when the pixel voltage reaches that of the detector substrate (zero bias). Technically, a pixel will continue to accumulate charge even into forward bias, but its response by that time will be significantly nonlinear.

Unlike the InSb detector, whose response is only weakly temperature-dependent, the NICMOS detector shows a significant temperature dependence of quantum efficiency and, possibly, related properties such as linearity and dark current. The IRIM detector is cooled with LN₂, at an approximate temperature of 75 K on Kitt Peak, but it is not absolutely controlled with a servo temperature controller. The actual detector temperature may thus depend on the ambient dome temperature and atmospheric pressure. Empirically, we have not seen significant differences in several linearity curves we have generated in the past, suggesting that such effects are small. Cautious observers may wish to take the time to generate a linearity relation during their run.

As the linearity relation indicates, while one may loosely define a full well of about 400,000 electrons, the response has already fallen by 6% at 350,000 electrons. The read noise and dark current of the array are quite low, as tabulated above. However, the dark current has a rather complex character and cannot be simply scaled with time. Necessary dark frames (e.g., for construction of flatfields or dark+bias subtraction of linearity sequences) must be taken separately at each integration time desired. In addition, it is strongly suggested that linearity sequences contain repeated "check" observations at one integration time (e.g., 1 sec) to verify the stability of the source. The array quantum efficiency improves by a factor of two from 60 K to 77 K, so we presently operate the device at the ~ 75 K temperature of LN₂ at Kitt Peak, and all tabulated values are referred to this temperature. There is a "charge retention" or "memory" effect whereby a small portion of incident signal from a given exposure survives many subsequent reset cycles, declining slowly. On stars the integrated value of the "memory" image is about 0.1% of the incident flux. A second aspect of this phenomenon is that the signal seen in a dark integration depends on how long the array has been exposed to sky.

Sky Background

Preparations

Object Coordinates for any epoch can be entered into the telescope computer for use during the run. This is often done by the telescope operator during the course of the night, but lengthy observing lists are best entered ahead of time by the observer (ask for help on this) or, preferably, by electronic submission (see below). These may include objects, standards, offset and guide stars, etc. Because of the large field of view, precise offsetting to an object is not critical. However, detection of very faint objects requires accurate long-term tracking combined with precision spatial modulation (dithering) to determine the sky level, a task for which the open-loop tracking of the telescope is usually inadequate. Use of the telescope guide probes for precision offsetting with reference to an off-axis guide star is highly recommended for registration of the many individual frames which such a limiting observation will require.

Conscientious observers may send coordinate lists via email (two weeks or more before the run) to coords@noao.edu. Files should be ASCII text, no longer than 2000 lines. Start the file with your name, a cache name, telescope, and dates of the observing run. Coordinates will be checked for format, loaded into the appropriate telescope computer, and acknowledgement will be sent. Each object should be one line of text. The format is object name, RA (starting column 16 or greater, delimited by first blank after col 15; hours, minutes, seconds), DEC (degrees, minutes, seconds), and epoch. Each field should be separated by one or more spaces (NO TABS); the delimiter in the RA and DEC fields may be spaces or colons. Example:

• alpha nuti 12:34:56.7 -89:59:59.9 1734.4

Photometric Performance and Flatfielding:

A series of experiments determined that exposing to about the same ADU level for each, a 10.5 and 14.0 K mag star scaled correctly to within 0.03 mag without linearity correction. Observing the same star at different locations around the chip and at different integration times gave rms scatter of a few hundredths of a magnitude after flatfielding, considered to be within the errors given the rather uncertain weather in which observations were made. Flatfield response using the sky is constant in a given filter over at least a factor of three in signal level. Flats in the J, H, and K' filters are very similar overall but show 5% differences when divided. Comparing dome flats to sky flats shows them equivalent to 1% at J and H but quite different at K', evidenced by a pronounced centro-symmetric pattern in the division. There is no quantitative information for the narrowband filters.

This is an IRIM-specific synopsis of the WILDFIRE manual written by Nick Buchholz. Observers interested in a more in-depth analysis of WILDFIRE are referred to that manual.

Initializing the Environment with OBSINIT

The optical CCD (ICE) and infrared (WILDFIRE) environments are both operated from the same account on the 2.1-m (2meter)and 4-m (4meter) telescopes. The all-important obsinit command performs a number of functions relevant to this operating procedure.

• Switching from ICE to WILDFIRE environment on the first night of an IR observing run.
• Efficiently and gracefully cleaning the disk of data from previous observers.
• Storing the current observer name and proposal number for archiving.

First Night of IRIM Block

On the first night of an IR block, the ICE environment will still be active (the presence of the "CCD Acquisition" and "CCD Reduction" windows will verify this). It will be necessary to run obsinit to change to the WILDFIRE environment, as well as for the other reasons above; since the hardware may be in an unknown state, it is recommended to run through a complete hardware initialization on the first night of an IR block as part of the obsinit process. This will involve rebooting the observer's SUN workstation with the DSP (in the computer room) powered on and the IRIM instrument power off.

• Verify that the DSP power is ON.
• Verify that the power to IRIM is OFF
• Quit both the "CCD Acquisition" and "CCD Reduction" windows
• Logout of any other IRAF processes
• Enter the command obsinit in a Shelltool or Xgterm window. One will be led through an interactive process:
  • name(s):
  • Proposal ID: (check the schedule or Preparation Form)
  • Operation (fire/ice): (enter fire)
  • Delete old data from disk and initialize (y/n): [this can take a while if you choose y]
  • Replace wfpar and tclSamples (y/n): (n will leave any changes)

Once this is complete, it is necessary to reboot the instrument computer with L1 A or Stop A; again, the instrument power must be off. After rebooting, the UNIX login prompt "[hostcomputer] login:" will appear; IGNORE THIS. After a few seconds, OpenWindows will automatically load and present the login window shown below:

Login as [telescope] with the current password posted on the workstation terminal. The WILDFIRE system will then load automatically, resulting in a terminal screen layout approximately like Fig. 4 below; the dashed window labeled Instrument Status will appear in the approximate position shown only after the instrument microcode has been loaded.

New Observer

On succeeding IRIM runs, obsinit is run only to enter the new observer and proposal ID information. It is NOT necessary to power down IRIM or reboot the computer. After logging out of all IRAF processes and running obsinit, simply exit OpenWindows from the desktop menu and log back in when the login window appears.

Normal WILDFIRE Startup

The Windows

Once the environment has been set to WILDFIRE by obsinit, it will remain in that state, even if it is necessary to reboot the instrument computer for any reason. There should be no reason to execute obsinit more than once during a run. If a reboot is required, the login procedure in the window displayed above will automatically bring up the WILDFIRE windows.

A brief description of the windows follows:

• 1. Instrument Control -- This is the window for entering all commands controlling the instrument or telescope. It will initially have a "[instrument computer]" prompt, and a "%" prompt when the instrument microcode is running.
• 2. IRAF XGTERM -- This window is used for IRAF commands for analysis of data or for shell commands such as creating directories, moving, or archiving data.
• 3. Instrument Status -- Once the instrument microcode has been loaded, this window will appear. Instrument, voltage, detector status commands in the Instrument Control window will output here. The existence of this window is a diagnostic of WILDFIRE, as a system crash will often close it.
• 4. Display -- This is an XIMTOOL window which may be accessed either through the Instrument Control window, where images may be automatically displayed as they are taken, or through the IRAF XTERM window via the IRAF display task. Those who are desperately attached to the SAOimage display may use it in place of XIMTOOL by killing the XIMTOOL process in the Console window (enter ps -ax in the Console window to get a listing of the process numbers, then kill -9 [process #]), and then entering saoimage & to bring up the SAOimage window.
• 5. Console -- This is generally used only for diagnostic or emergency purposes, such as killing a hung process or the entire WILDFIRE when it hangs up.

Bringing up WILDFIRE

There are three basic steps in the complete startup of WILDFIRE: hardware initialization; starting WILDFIRE; instrument initialization. The procedure below will go through all three steps, as would be necessary on the first night the instrument is on the telescope.

Hardware Initialization

This procedure establishes the link between the DSP box and the computer, by rebooting the observer's SUN workstation with the IRIM power off. The obsinit procedure for the first night of an IRIM block (described above) includes these steps.

Starting WILDFIRE

The installation of the instrument and cables will be handled before the beginning of the run by the mountain technical staff and are not of concern to the user. IRIM remains on the telescope for the entire observing run and the LN₂ cryogen flasks are filled twice per day by the observing technicians.

Cable Harness -- Three cables run from the telescope junction box to the IRIM DCU. Two of these are power for the analog and digital electronics, the other a fiber optic cable containing six individual fibers. Four of these, marked TRANS1, RCVR1, TRANS2, RCVR2 provide communication to the DSP in the computer room, and are connected to the appropriate connectors on the DCU. The other two are spares for use in the event of a failure. Caution: The fiber optics are delicate and should not be subjected to force or bending of short (< 7 cm) radius; they should be secured to the instrument to relieve any strain on the connectors.

Dewar Cables -- There are two cables from the DCU box to the rest of the instrument:

• ACU Power: A single cable from the DCU to the ACU provides analog power.
• Filter: This is a Y-cable from the DCU to the filter drive motor and encoder.

In addition, four ribbon cables carry commands and data from the DCU to the ACU. These should always be left in place.

Getting Started

After IRIM is installed on the telescope, go through the WILDFIRE startup procedure outlined previously. Once the system is operational and the detector activated, check the detector and temperature status with status s and compare with the nominal values below:

Verify that the filter wheel is functional by moving to a number of filters using the command fwl to [filtername]. The successful completion of a motion should return a verifying message. The command fwl ?pos should return the current filter position.

Dark Current and Noise

After the detector is stabilized at operation temperature, one may check the dark current and read noise with a series of dark bias observations. With the filter wheel at the "dark" position (fwl to dark), take a large number (10 - 50) of observations at the minimum integration time of 0.38 s. Using the IRAF 'imcomb' task (mode=average, sigclip), calculate the average and sigma images. Examine the statistics of the average and sigma images with 'imstat', using a [50:200,50:200] subarray to avoid the bad regions at the edge of the array. If IRIM is operating properly, one should obtain approximately:

Mean Image : mean ~ 4.5 ADU

Sigma Image: mean ~ 3.2 ADU

If the value of the mean image is much greater than 5 ADU, there is additional dark current, most likely a result of the system not being completely cooled down. A sigma image value much greater than 3 ADU (30 e-) suggests high read noise and a call for help is in order.

Twist & Shout

This one-time exercise aligns the optical axes of the telescope and instrument, by the adjustment of a gimbal mount within the rotator. The alignment is determined by use of the "Christmas Tree" lights, which mount at the periphery of the secondary mirror at the N,E,S,W directions. Using a narrowband filter, one checks the brightness of the image for opposite lights and adjusts the tilt in that direction until the illumination is equal; this process is then repeated for the other coordinate. The primary covers should be open just enough to view the secondary. This will be done during the first night of setup and should not need repeating during a blocked run.

Techniques

Focus

IRIM is not fully achromatized, and there will be focus differences between filters. At the beginning of a run, one should determine these offsets empirically for all filters one intends to use. To focus the telescope in a given filter, use a star which gives 5000-10000 ADU peak signal in a few seconds' integration. Initial focus is easily done by observing using movie (continuous integration and display), changing the telescope focus and noting the character of the image vs focus readout. For fine adjustment, take a series of frames covering several focus steps across what appears to be the best value. Use the IRAF 'imexam' task to judge the best setting. It is dangerous to query the display with 'imexam' while running in movie mode. If IRAF and instrument tasks access the display window simultaneously, WILDFIRE will crash.

A recent determination of the focus settings at the 2.1-m and 4-m telescopes in the J, H, K, and K' filters yielded the following focus readout offsets relative to J (using the absolute value of the readout in both cases):

Although this table should provide a rough guide for the focus offsets, we still recommend that they be determined empirically during an observing run.

Changing temperature of the telescope structure will also require refocus. A recent determination of focus vs temperature at the 4-m is tabulated here, courtesy of S. Courteau and J. Holtzman.

Keep in mind that the zero point of this relation will change from run to run as IRIM is remounted on the telescope, although the slope of ~37 units/ °C should be the same.

Observing

Because of the high brightness of the infrared sky, often greater than that of the objects under study, and the complex dark current, the technique of subtracting a bias (dark) frame and dividing by a flatfield works poorly in the infrared. Some technique of subtracting the sky prior to flatfielding is necessary. In practice, this is accomplished by several observations of the object field, separated by small motions of the telescope, so that the stars in the field are imaged onto different sets of pixels in each of the images (this is sometimes referred to as "dithering"). If the spatial object density of the field is relatively low, then combining the images using a median algorithm will result in an average from which the astronomical sources have been removed; i.e., a sky frame. This may be done using the IRAF 'imcombine' task. If the sky amplitude varies from frame to frame, it may be necessary to scale the frames by the mean to obtain a good median average. The resultant sky frame may then be subtracted from each of the raw images to yield sky-subtracted images which then may be flatfielded and reduced. Note that the detector dark current is also removed in this process. Creating the sky frame from a number of raw images also reduces the increase in statistical noise resulting from arithmetic operations on the images.

If the object field is extremely crowded (e.g., a globular cluster) or contains an extended source, it is necessary to move the telescope a significantly larger distance ("wobbling" or "jogging") to a suitably sparsely populated "sky" field. However, succeeding observations within both the object and sky fields should still be "dithered" so that one may obtain an average sky frame using the median averaging technique and to avoid imaging the same parts of the object field onto defective pixels. As noted below, there are some instances in which this method introduces other problems.

There are a number of resident scripts which will automatically carry out these types of observations. In addition, existing scripts may be copied and customized by the observer to carry out more complex or exotic observations.

Flatfielding

Because sky flats provide the same array illumination as real observations, they are preferable in principle to dome flats using the White Spot. It is, nonetheless, a good idea to obtain dome flats as a backup. If one is observing in a sufficiently sparse star field, one may use the same set of observations for the object, sky, and flatfield.

Because the sky flats will include the array dark current, it is necessary to obtain separate "dark" observations for subtraction from the sky observations. Unfortunately, what constitutes a "dark" frame for creating flats is ill-determined due to the "memory" effect which accumulates in time. For example, a series of observations in the dark filter immediately following sky (or dome) flat observations will show a monotonic decrease in mean value as the "memory" of the relatively bright preceding observations slowly decays. By the same token, a series of dome flats following dark or low-background observations will show a monotonic increase in mean value as the "memory" of the higher flux observations accumulates. One possible approach is to take a larger number of dome flat or dark observations and reject those early in the series, when the change in value from one frame to the next is the greatest.

Photometric Standards