Return to TIFKAM Index
The MDM/Ohio State Array Infrared Camera (MOSAIC) is a multi-purpose infrared
imager and spectrometer. The instrument was built as collaboration between
MDM and Ohio State and was designed to work at the f/7.5 focus of
the Hiltner 2.4-m telescope at 1-2.5 micron wavelengths. The instrument
currently uses a 512x1024 InSb array supplied by NOAO. Internal optics
allow for two plate scales and a variety of spectroscopic resolutions.
Internal mechanisms control the selected filter, focal plane mask, pupil
mask alignment, and optical arrangement. All mechanisms are user controlled
and can be used to rapidly reconfigure the instrument, which allows for
substantial observing flexibility.
Because MOSAIC is also the name of the 8K X 8K CCD imager at KPNO, the present
instrument is referred to as TIFKAM or ONIS (Ohio State/NOAO Imaging Spectrograph)
while in use at KPNO.
TIFKAM is a general-purpose 1-2.5 micron imager and moderate-resolution
spectrometer built as collaboration between Ohio State and the MDM
(Michigan-Dartmouth-MIT)
Observatory. The instrument was designed to fully illuminate a 1024x1024
InSb array, but currently has a 512x1024 InSb array (supplied by NOAO and
the USNO). The instrument can be used at any telescope with a beam slower
than f/7.5 and includes cameras for two plate scales (approximately
f/7.5 and f/16.5) and pupil viewing. Grisms provide spectroscopic
dispersion with resolutions of ~600 to ~1200 (assuming a two-pixel slit)
depending on the camera and grism used. All configurations of the instrument
are under user control and can be selected rapidly (i.e. within a few minutes),
so many different observational programs can be accommodated during a single
night.
TIFKAM was designed to work at the f/7.5 cassegrain focus of
the MDM 2.4-m Hiltner telescope, but has also been used at the MDM 1.3-m
and KPNO 2.1-m and 4-m telescopes. In exchange for the large detector array,
the instrument is available for community use at KPNO.
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 control interface for the instrument is the Prospero package
written at OSU. More information on Prospero can be found on the
Web at
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, copies of which
may be obtained from:
Or via our anonymous ftp server (
ftp.astronomy.ohio-state.edu,
directory pub/prospero/instcards). A separate document, the
Prospero Observer's Guide, describes how
to use the Prospero package.
Pictures of the instrument can be seen here
along with a gallery of sample astronomical data on associated pages.
The instrument has many operational modes and has not been extensively
tested in each. In particular, we have had very little photometric weather
(thanks to El Nino) and have measured the throughput in the various
spectroscopic modes only on the KPNO 2.1-m telescope.
There are 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.
Feb 1999. The Ks filter previously in TIFKAM has been replaced. The
new filter has significantly higher throughput, yields good images, and
does not require a blocking filter.
TIFKAM has two camera lens sets (f/7 and f/16.5). The
table below summarizes the possible plate scales:
Here are representative photometric results.
The band gap energy of InSb is ~5 times smaller than that of Si, so the detector
array is sensitive to radiation out to wavelengths around 5500 nm. The
relatively small band gap energy also causes high dark current at relatively
high temperatures. If this array is operated at 77 K,
for example, the dark current would be ~1X106 e-/sec.
Consequently, the array must be operated at a temperature significantly
colder than LN2. TIFKAM has a small mechanical refrigerator
(a Split-Sterling Cycle cryo-cooler) that cools the detector and its mount
to ~35 K.
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. This task is handled by the OT, usually prior
to opening up in the evening. The OT has been requested to log each fill
of LN2, along with the ambient temperature, detector temperature
LCD readout, and power setting of the cryocooler into the logbook which is
kept with the instrument.
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.
We do not know if the cause of the residual is our particular read-out
electronics (reset voltages, bias level, etc.) or if it is intrinsic to
the detector array. Other InSb arrays operated by KPNO do not seem to show
this level of residual, but the particular detector in TIFKAM was not carefully
evaluated using a KPNO system.
We suggest that if a 1% residual image will seriously compromise your
results, you dither the telescope faithfully and, perhaps, read the array
seveal times (set the exposure time to 0 seconds, do an mgo 3,
and throw those images away) between each science exposure.
One may use the NOAO Exposure
Time Calculator to estimate the performance of TIFKAM for imaging
programs. However, it is very important to understand the assumptions and
limitations of this tool, and we strongly suggest reading the background
material on Signal-to-Noise Calculations.
Note:When answering questions such as the desired data directory,
it is necessary to enter the value (e.g., /data1/4meter/data) explicitly.
Even though the default value is given in square brackets, entering
a carriage return is not sufficient!
You should be presented with two windows: an xterm labeled
"Prospero Command", and a second window labeled "Prospero
Status" located immediately above the command window. The status
window will be mostly blank until you have connected Prospero to
the data-taking PCs and started an observing session.
At KPNO, caliban is supposed to be started automatically
upon login to the workstation after completing the osuinit process.
If caliban is not running (or has crashed), it needs to be restarted by
typing in an xgterm window on the workstation
This appends the ".fits" extension to FITS files created by the system
to enable reading them with the FITS kernel which
became a standard feature with version 2.11 of IRAF in summer 1997.
The caliban setup at KPNO is also responsible for archiving all
data taken with the instrument using the observatory's Save-the-Bits
Archive. Archive
Status Reports are available on-line via the Web.
Caliban will start up as an icon labeled
"Caliban" and automatically connect to the WC and mount the two
data transfer disks. Normally caliban is left as an icon to save
on workstation screen space, and is only opened if you need to interact
directly with caliban to investigate problems. Typing the ? in the
caliban window will give a list of interactive commands. Be aware that careless
use of the caliban command window could shut down the data transfer system
and lead to loss of data (if you don't have a good reason to mess with
it, you probably shouldn't).
Once ariel and caliban are running, go to the Prospero command
window and enter
There will be a bunch of chatter between Prospero and the various
processes, and you will see the Prospero Status window begin to
fill up with detector, instrument, and system status information.
A detailed startup example is given in Appendix C of the
Prospero Observer's Guide, along with discussion of possible
error situations.
Finally, in the Prospero command window, enter
This will set up specific information for a night's run, such as the observers,
proposal ID, data directory, filename, etc. It is a good idea to execute
runinit at the beginning of each night, as it preserves various
state variables and permits a quick recovery in case of a crash. In addition,
the IC and WC configurations will also be saved for the same reason.
1. Introduction
2. System Overview
3. System Characteristics and Performance
Imaging
TIFKAM was designed for good image quality over the entire field of a 1024x1024
detector array. The optical design has a relatively small number of surfaces,
so the instrumental throughput should be quite good. The terrible weather
conditions during late 1997 and early 1998 has delayed accurate measurement
of the actual throughput. One set of measurements at the KPNO 2.1-m suggest
a total throughput (including the atmosphere and telescope optics) of approximately
0.25 in the J and H filters and 0.30 at K. More recent measurements in
June 1998 at the KPNO 4-m under excellent conditions yield values 0.03
to 0.05 higher.
MDM 1.3m
...
...
...
...
...
...
...
...
KPNO 2.1m
61500
56500
30000
37300
30
130
250
125
MDM 2.4m
83000
63000
33000
...
...
...
...
...
KPNO 4m
220000
196000
100000
...
...
...
...
...
Spectroscopy
Spectroscopic capability is provided 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. 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 2400 nm due to the transmission of the filter);
we are currently trying to obtain special filters to use specifically as
spectroscopic blockers. The table gives the resolution assuming the 50
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 or when observing spectroscopically on the KPNO 4-m (where the 50
micron slit is only 0.36" wide).
grism
blocking filter
resolution (2 pixels)
full array wavelength coverage (nm)
illuminated rows
J/K
J
1450
1110 - 1360 (2nd order)
400 - 982
J/K
K
1325
1970 - 2420 (1st order)
160 - 704
H
H
1150
1480 - 1800 (2nd order)
470 - 918
J+H
IJH
720
950 - 1910 (1st order)
1 - 900
JHK
JHK
750
1220 - 2490 (usable to 2290; limited by blue leak)
1 - 1024
Signal Levels
In order to estimate the spectroscopic throughput, we obtained slitless
spectra of the Elias standard HD84800 in April 1998 at the KPNO 2.1-m
telescope under clear but rather poor (1.5-2 arcsec) seeing conditions. The
latter factor is not important in the throughput measurement, but
it made an estimate of the typical slit losses impractical. The
three primary configurations (J filter, J/K Grism; H filter, H/L
Grism; K filter, J/K Grism) all yielded approximately the same
maximum signal (~5 X 105 ADU/s) referenced to a 0.0
mag star and an integration time of 1 s. These are the signals
integrated through the spatial profile. The peak signal on a given
row, which is an important consideration with regard to possible saturation, will
naturally depend on the seeing. These spectra have been bias subtracted,
but not flatfielded, as they are intended to be a guide to the
raw signal level.
Sky Background Levels
On the same night, we measured the sky background through the 4 pixel
slit, using the same three configurations. An equivalent dark
exposure was used to subtract the bias, but the baseline subtraction
may not be perfect. The plots, however, give a representation of
the expected background. Use of a narrower slit will decrease the
continuum in the K band proportionally, but will not affect the
height of the OH emission lines, as they are unresolved.
4. The InSb Detector Array
Introduction
The infrared array currently in TIFKAM is a 512 X 1024 InSb array on loan
from NOAO/USNO. The array is a product of the NOAO/USNO Aladdin
development project and is one of the largest format infrared detectors
in routine use. The array is a hybrid of a silicon multiplexer and an array
of infrared sensitive detectors. The two pieces are pressed together, with
indium bumps on each piece making the electrical connection. The multiplexor
is an array of discrete read-out transistors and is unlike an optical CCD.
In particular, the array can be read non-destructively and each pixel can
be read separately and in any order. Although these differences do not
matter for typical observing programs, future upgrades could permit additional
observational flexibility. For example, designated sub-arrays could be
read-out at high frequency, then shifted and added later to create diffraction
limited images.
Pixel Size
27 microns Operating Temperature
35 K Detector bias
300 mV System gain
4.0 electrons/ADU Read noise
35 electrons RMS Dark current
Approximately 1.2 electrons/s Full well capacity
Approx 105 electrons Minimum integration time
1.7 s
Dark Current and Read Noise
The dark current in the array is low. We currently measure ~0.3 ADU/sec
dark current, some of which may be due to background radiation in the dewar.
The dark current can be higher if the array cryo-cooler has lost power
(the array is warming up) or if the instrument has been cold for less than
~36 hours. The
temperature of the array is given on an LCD on the side
of the large aluminum box (the unanodized, silver one). At the nominal
detector temperature of 35 K, the LCD should read 4.5, with a slope of about
0.3/K. Values between 4.0 and 5.0 indicate a normal detector temperature.
If you measure dark current
significantly higher than 1 ADU/sec, please check the LCD and report the
problem to Dick Joyce (rjoyce@noao.edu),
Mike Merrill (mmerrill@noao.edu), or
Darren DePoy (
depoy@astronomy.ohio-state.edu).
The read noise of the array is ~9 ADU. The effective read noise of the
detector can be reduced by reading each pixel's signal non-destructively
many times; factors of 2-3 improvement are possible. Note, however, that
a read noise of 9 ADU is low compared to the shot noise for signals greater
than ~1000 ADU, which happens quickly in most imaging and low-resolution
spectroscopic observations. If your observations are read-noise limited,
please contact Darren DePoy (
depoy@astronomy.ohio-state.edu)
for more information.
Maintenance
Linearity and Saturation
The array becomes significantly non-linear before the 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
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%.
Bad Pixels
There are many dead, hot, unresponsive, and just plain bad pixels on the
array. There is a large group of essentially dead pixels in the upper left
corner of the array (referred to as the "fingerprint" region for obvious
morphological reasons) and many scattered bad pixels in the upper right
corner of the detector. The large number of bad pixels makes it important
to "dither" images so that all pieces of the sky are measured by a good
pixel. We have a bad pixel mask dating from
November 1998 stored here for convenience, although it is possible that
this may become outdated in time.
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.
System Sensitivity
5. Observing Setup
Unix Workstation Startup
Details of how to use the Prospero program for observing are described
in detail in the Prospero Observer's Guide (available in the control
room or on-line from the OSU anonymous ftp site, see
section 8 below).
This section describes the setup procedures you need to follow on the first
day of your observing run (or to do a full-dress recovery after a major
system crash).
KPNO Startup
At the very beginning of your run at Kitt Peak, you need to configure the
observing account for running the Prospero data-taking package,
and delete the previous observer's data. This is accomplished for you by
running the osuinit script. The procedure is as follows:
Among other things, osuinit will delete the previous observer's data files
and reset the data storage directory used by the data-taking system for
your images. You have to make sure that IRAF is closed and that
the ariel and caliban daemons are not running before starting
the osuinit script, as these will not be correctly reset if left
running while osuinit is executing.
Starting Prospero
At Kitt Peak, Prospero should startup automatically when you log
back in after running the osuinit script. If you need to restart
Prospero during an observing run, use the workstation menu (left
mouse button) and select the option labeled
Data Acquisition (Prospero) from the menu.
Other programs needed on startup
Before Prospero can be connected to the data-taking PCs, you need
to make sure that two other support programs are running: ariel
and caliban.
ariel
The ariel program mediates communications between Prospero
and the PC computers running the ICIMACS system.
Under ideal circumstances, ariel should be executed automatically
when the WC is started up, although in practice, it is usually
necessary to manually start the ariel program by entering
ariel at the keyboard of the WC.
caliban
The caliban program is responsible for transferring images in FITS
format from the data-taking PCs to the Sparcstation. Before starting caliban,
the ICIMACS system has to be turned on and the WC computer started
and connected.
prospero
Setup Checklist
Pupil Mask Alignment
There is a silicon lens in the instrument that focuses the telescope pupil
onto the detector array. There is also a cold pupil mask in the dewar that
can be moved (using the xpupil and ypupil commands)
to align the mask
with the pupil to reject the thermal emission from the telescope secondary
support and mirror cell. The procedure for aligning the mask is to insert
the silicon lens and the K filter into the beam (see the list of wheel
populations for the instrument to determine
which turret and filter position is appropriate), set the camera focus
to 2000 (camfocus 2000), and then issue the ie campupil command.
The ie campupil is a direct command to the instrument electronics
(hence the "ie" prefix) which
offsets the camera turret to a position where the entirepupil
can be seen on the detector array (note that this is not a legitimate location
for the camera turret, but is between detented camera locations).
A short exposure against the night sky (clouds will do, but the interior of the dome is probably too bright) ideally should look like the figure below. The bright areas are the telescope, secondary, support structure, etc. that should be masked off by the cold stop. Adjust the values of xpupil and ypupil until the dark areas obscure as much of the bright as possible (the current mask is a compromise choice; it is not optimal for any particular telescope but includes a large central obscuration for the KPNO 2.1-m, an outer diameter for the MDM 2.4-m and KPNO 2.1-m, and secondary support vanes for MDM telescopes). Note that at the KPNO telescopes, the secondary support vanes are not aligned with the mask vanes and will always be bright. Because the KPNO 4-m primary is masked down, the TIFKAM pupil stop is somewhat oversize and one will see a bright rim to the pupil image.
Note that this procedure need only be done once per run.
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. Steps of 50-100 in the camera focus are appropriate. Approximate values (March 1998) are camfocus 100 for the f/7.5 camera, camfocus 1100 for the f/16.5 camera, and camfocus 2000 for the pupil viewing silicon lens. The best focus for the coronagraphic mask was determined to be approximately camfocus 350, but we recommend that potential users of this feature run through a focus sequence to verify this value. For spectroscopy, we recommend that the optimum camera focus be determined using a calibration line source. We found that the JHK grism yielded the optimum resolution at camfocus 250, even though the slit was focused at camfocus 100 in direct imaging. 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.
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. There is no analog to the CCD command that reads out a single frame with multiple images of a star at various telescope foci. At the KPNO 2.1-m, the f/8 secondary readout at focus is approximately 15000.
Note that all the filters and prefilters are in collimated light, so there is no change of optimal focus with wavelength.
The syntax for moving these mechanisms is of the form [mechanism] [value]
. To set up for J band imaging with the PK-50 blocker, for example,
one would execute prefilt 2 and then filter 6.
The plate scale at the KPNO 2.1-m was measured to be 0.341" per pixel
in both the J and K filters by observations of NGC 7790 in September
1996, equivalent to a total field of view 175 X 349 arcsec. The
scale at the 4-m telescope was measured to be 0.178" per pixel
in June 1998. Based on the signals observed at the 2.1-m,
the total (instrument + telescope) throughput is approximately 0.25 at J
and H (both of which require a PK-50 blocker) and 0.30 at K. The slits are roughly
480 pixels long. The 2 pixel slit is slightly rotated with respect to
the array, by about 2 pixels over the length of the slit. The 4 pixel
slit appears aligned with respect to the array to within 1 pixel.
There are three ways of taking exposures: go, mgo [n],
and avego [n]. go takes a single exposure;
mgo [n] takes n integrations and writes n frames
to the disk; avego [n] takes n integrations and writes 1
averaged frame to the
disk. The array becomes significantly non-linear beyond ~8000 ADU, so
it is advisable to keep the integration times sufficiently short to keep
the background less than this level. At this point
the background shot 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.
Mike Merrill's provisional IRAF script tmove may be used for centering
stars on the array, using an image displayed in the ximtool window. Because
this is not yet a standard IRAF task, it will probably have to be
manually installed for an observing run.
task tmove = /data1/4meter/tmove.cl
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
source such as a comparison lamp or the atmospheric OH lines.
There is a fairly strong focus dependence on temperature at both the 2.1-m
and 4-m telescopes; in addition, the 2.1-m shows a pronounced effect
with zenith distance. Once you have a good focus, the coefficients below will
guide you in refocusing as the temperature and airmass changes.
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.
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.
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.
To download the complete document package (compressed tar archive):
The TIFKAM support scientists for KPNO are Mike Merrill mmerrill@noao.edu and Dick Joyce rjoyce@noao.edu.
6. Observing Techniques
Imaging Observations
There are 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., which then lists out the populations of that particular mechanism.
Tmove
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 2.1-m and 4-m
telescopes.
Telescope Focus
Focus Change ( 
4m
-90
0 2.1m
-75
-130 Prospero Scripts
7. 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.
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
provide a useful grid for spectral calibration. However, at both the
2.1-m and 4-m telescopes, the spectral comparison sources built into
the guiders provide an excellent Ar spectrum, with a few He and Ne
lines thrown in. One may obtain a high S/N calibration spectrum
with a 30 - 60 sec exposure (followed by an equal dark observation
for bias subtraction) using the HeNeAr source, so this technique
is recommended. Plots of HeNeAr spectra for the five presently
available configurations are given below. These were obtained
on the KPNO 2.1-m telescope with the 2 pixel slit.
Flat Fielding
The InSb detector array is reasonably flat; we have measured ~2% photometric
accuracy for standard stars (at K) without flat-fielding the data. At the
KPNO telescopes, 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.
8. Additional Information
World-Wide Web Services:
MOSAIC/TIFKAM Web Page:
Prospero Support Web Page:
Documents available via Anonymous FTP:
Connection info:
MOSAIC/TIFKAM Documents:
PostScript format copies of all MOSAIC/TIFKAM documents are available online
at our anonymous ftp site as well as through our WWW sites above:
cd pub/MOSAIC
Documents are available as either gzip compressed PostScript (.gz) files
or uncompressed PostScript (.ps) files. For example:
binary
get mosaic.gz
or
ascii
get mosaic.ps
and print using
zcat mosaic.gz | lpr
for gzip compressed files.
binary
get mosdocs.tar.Z
and unpack using
zcat mosdocs.tar.Z | tar xvf -
See the README file for the list of available documents.
Prospero Documents:
PostScript format copies of the various Prospero manuals and reference
guides are available online at our anonymous ftp site as well as through
our WWW sites above:
cd pub/prospero
Individual files are available as either gzip compressed PostScript (.gz)
files or uncompressed PostScript (.ps) files. For example, to retrieve
the observer's guide, you would type either:
binary
get obsguide.gz
or
ascii
get obsguide.ps
To retrieve all documents (Unix tar file):
binary
get prdocs.tar
and unpack using
tar xvf prdocs.tar
A gzip-compressed archive is available as prdocs.tgz. See the README
file for the list of available documents.
KPNO Support Staff
14 Sep 1999
