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| SYSTEM LEVEL | ||
|---|---|---|
| ! commands ? command ed name help name man name source program |
execute the commands in csh or run csh
give help on a command call up an editor on a proc display help for one of these topics display a man page for a given topic temporarily include tcl program within recognized system; need to source again after powerup or go (full path name required) | |
| DETECTOR | ||
| activate deactivate setup crsp |
activate the detector deactivate the detector set up the default CRSP voltages and prompt for activation | |
| WILDFIRE | ||
| startwf exit trouble hung |
initiate bootstrapping and downloading of the WILDFIRE system
deactivate the array and exit the WILDFIRE controller open troubleshooting session (do NOT enter in Instrument Control window) attempts to complete link protocol; used as part of the restart procedure when WILDFIRE is hung (INSTRUMENT CONTROL window unresponsive and data collection stalled); must be entered in Console window | |
| HOUSEKEEPING | ||
| status [|s|v|t|f] longheaders [on|off] tcp_[on|off] |
display a status screen; (general status |s|;
voltages|v|; temperatures|t|; filters |f|)
will disable/enable house keeping data in the header enable/disable link to TCP for telescope status info and offsetting | |
| PARAMETER FILES | ||
| Note: a parameter has two attributes, its value and flags indicating whether the parameter and its value should be displayed and/or queried when the ask or observe tasks are run. | ||
| lpar plist psave filename puse filename ped ask eask |
list the names of the available parameter files list all the current parameters save the current parameter set (values and ask/display flags) to the named parameter file load the named parameter file edit the current parameter file selected by 'puse' asking all questions without regard to their query status prompt for the 'eask' selected subset of parameters within the current parameter set iterate through all the known parameters, allowing the user to specify which parameters are queried and which are displayed. After each question an "l" signifies display only; "a" signifies query; "la" will list the current value (which may be selected by [cr]) or accept a new entry | |
| ACCESSING INDIVIDUAL PARAMETERS | ||
| ?coadds coadds [n] ?lnrs lnrs [n] pics [n] set-time [f] nextpic [n] header_dir pixel_dir mode title comment set fdly [f] resetoffset filename filename |
returns the number of coadds for next image set the number of coadds for next image; if no argument given, will prompt with current value returns the number of low noise reads per coadd sets the number of low noise reads per coadd; must be <=16! sets number of pictures to be taken at each observe/go. When no argument given, prompts with current value sets the integration time (to millisec level); when no argument given, prompts with current value sets the picture index appended to filename to [n]; when no argument given, prompts with current value sets IRAF path for image headers sets IRAF path for pixel files sets operational mode for array readout (stare, sep, hphot) sets title field for IRAF image header sets comment line within IRAF image header set the reset-read delay time to [f] seconds reset the "offset" values in the header to 0,0 sets the IRAF image "filename". The path is is not included in "filename"; if no argument given, will prompt with current value. For CRSP, a "%d" or "%03d" should be inserted where the picture number should be placed. If no field is given, "%03d" will be appended. The format will be: "filename"//"nextpic" WARNING!! SAVED IMAGES WILL OVERWRITE EXISTING IMAGES IF THE FULL FILENAMES CONFLICT!! | |
| OBSERVING | ||
| observe go abort save only filt to [n] slit to [n] grat to [n] chgrat lambda [m] [f] rot to [n] sangle [f] ?filter ?sangle ?lambda east [n] west [n] north [n] south [n] toffset [e] [n] zs [z1] [z2] zs 0 0 movie |
perform one observation using current parameter set,
prompting for key parameters initiate an observation using the previously set parameters abort an observation (enter in Instrument Control window) enter immediately following an abort to ensure that future images are saved to disk position filter wheel to filter [n] by name or number position slit wheel to slit [n] by name or number move the grating to encoder value [n] move grating to 4000 ecu for grating change move grating to wavelength [f] in order [m] move rotator to encoder position [n]. The angle [pa]=([n]-334) / 20.436 move rotator to position angle [f] on sky displays position of grating, filter, slit, rotator, selected grating, and any limit errors displays current rotator position angle query order [m] and wavelength [f] for the current grating position move telescope [n] arcseconds east move telescope [n] arcseconds west move telescope [n] arcseconds north move telescope [n] arcseconds south move telescope [e] arcsecs east and [n] arcsecs north: + for north/east; - for south/west set zscale values [z1] and [z2] for the image display enables autoscaling for the image display begin observe/display loop. NOTE: parameters (filename, running number, integration time, coadds, etc.) will be those of previous observation or 'ask' routine unless specifically reset!!!! Movie frames are saved to disk and should be deleted periodically. It helps to use a filename like "junk" when using movie. Terminate movie with end [CR] in Instrument Control window. | |
| no signal | Mirror covers, dark slide closed. If stars visible on TV, check ?filter for proper filter, slit, grating position. Check status s for proper temperatures, voltages. Check that green LED in analog electronics box is lit. If the detector has been accidently deactivated, program will not sense this; observe will work, but return pixel values near zero in image. | |
| apparent vignetting | Filter may have not seated properly in detent, resulting in vignetting of one side of slit. Cycle through the filters and return to the desired filter. Take a movie image sufficiently long to measure sky continuum or lines; both images of the slit should be sharply defined. | |
| poor hold time | He gas in vacuum jacket. If He hold time is low and boiloff from vent line is high, or LHe temperature sensor reads lower than 3.0, it will be necessary to pump on the vacuum jacket with the portable Tribodyn pump. Call for technical assistance. | |
| rotator slips | Check for cables snagging instrument. At extreme zenith angles, imbalance may overcome slip clutch. One may manually assist rotator to desired position angle. Slippage may result in message "servo 2 failed -- JAMMED", and it will be necessary to repeat the rotation command. | |
| motors inoperative | If only one or two motors will not work, check ?filter for error message. Check cables on instrument (Figures 3, 4, 5 in Manual). If ?filter gives short or long limit error message, it will be necessary to move the grating out of the limit manually. Call for assistance. | |
| "JAMMED" motors | The motor code checks the motor encoder positions periodically during a motor motion. A discrepancy in the actual and calculated encoder readings will stop the motion and return a message "servo [n] failed -- JAMMED". This does not necessarily indicate physical jamming. Repeating the motor command will often complete the motion successfully. | |
| bootstrap failure | If the startwf procedure fails during the "bootstrapping node ..." process, a likely culprit is a bad (or incorrect, if the failure occurs on the initial setup) fiber optic connection. Check the three status LEDs visible through sight holes on the DCU. The top LED should be green if there is power to the instrument; the two lower LEDs should be off. If either the middle (channel 1) or lower (channel 2) LED is red, there is a fiber continuity problem in that channel. There is a duplicate set of LEDs in the Heurikon DSP box in the computer room; it is necessary to remove the front cover to view them. A bad fiber channel will require the substitution of one of the spare fibers. Call for assistance. |
The following procedures are intended as a guide for restoring the WILDFIRE system following various levels of system failure. Re-booting the computer and cycling power to the instrument or DSP in the Heurikon box in the computer room are not normal WILDFIRE operations and should not be done without proper consultation, or unless the specific conditions below are valid. These procedures are listed roughly in order of increasing severity, so unless a specific condition has occurred (e.g., DSP power cycling), try the less dramatic procedures first.
An extensive troubleshooting library may be consulted by entering trouble in any active window (except the Instrument Control window). The resulting interactive session can be used to diagnose and correct problems.
If the Instrument Status window has vanished, first check to see if it has simply been closed. Type fireproc (or !fireproc if necessary) from an active window and look for the "hkserv" process.
If the process is present, the window has been closed, and it will be necessary to locate and open it. If the icon is not visible, it may be hiding behind one of the open windows. In OpenWindows, one can check the "windows" item in the menu for the status of all operating windows; if the Instrument Status window is present, open it and continue observing.
If the Instrument Status window has died, perform the SIMPLE RESTART procedure below.
If WILDFIRE has crashed (Instrument Status window has vanished and could not be found by above procedures), and/or the "[hostcomputer]" prompt has returned to the Instrument Control window, the following steps within the Instrument Control window should restore operation:
[NOTE: If the power to the instrument and/or the Heurikon DSP box in the computer room has been interrupted or the computer has been rebooted, this procedure may not be sufficient. See below for more specific procedures]
If WILDFIRE is hung (Instrument Control window unresponsive and data collection stalled):
If the STALLED SYSTEM procedure fails to return the UNIX prompt, or an examination of the operating processes by entering ps ax in the Console window reveals a process which cannot be halted via the kill -9 [process number] command, it will be necessary to reboot the instrument computer.
If the power to the instrument was interrupted but the black Heurikon DSP box in the computer room remained powered up and the computer was not rebooted:
If the black Heurikon DSP box in the computer room has been powered down, then it is necessary to do the following. NOTE: THE ORDER OF THESE STEPS IS IMPORTANT. IF THE DSP BOX IS POWERED DOWN, REBOOTING THE INSTRUMENT COMPUTER IS NECESSARY. MAKE SURE NO ONE ELSE IS USING THE INSTRUMENT COMPUTER AT THE TIME. IF ONLY THE POWER TO THE INSTRUMENT HAS BEEN INTERRUPTED, PERFORM THE PROCEDURE ABOVE. WHENEVER THE INSTRUMENT COMPUTER IS REBOOTED, THE INSTRUMENT POWER MUST BE OFF AND THE HEURIKON DSP POWER ON!
Below are some recommended grating settings in ECU for CRSP.
The low-resolution grating 2 easily covers a given spectral window at one setting, and
the intermediate resolution grating 3 can cover
the entire J band and a significant fraction (2.01 - 2.42 micron) of
the K band at one setting. Grating 4 provides high efficiency intermediate
resolution in the I, H, and L bands, complementing grating 3.
The settings in the tables below are designed to cover the spectral windows
in the smallest number of grating settings.
The settings recommended for grating 2 approximately
center the window on the array, except for the L band, where it
is chosen to put the long end of the band at the edge of the
array, to prevent overloading the array with background. The wavelength given is
for the center of the array (row 128). The tables include the following
zero-order offsets:
Wavelengths determined from these tables should be accurate within a few
pixels. For precision wavelength determination, it will be necessary to
observe either a source with known spectral lines or a calibration lamp,
or utilize the atmospheric airglow lines as a calibration grid.
[Note that wavelength increases from bottom to top on the array.]
Note: Should one wish to cover a different spectral range than the standard
windows given, it is best to use the lambda command to center an
explicit wavelength on the array.
For initial focusing at zero order, set grating 1 to 128 ECU. This may
be accomplished for any grating with the command lambda 0 0.
The grating change setting is 4000 ECU. The command chgrat
will execute a motion to this position.
Note that the wavelength increases from bottom to top on the array and that
the dispersion is a function of the grating angle, so reduction of each segment
of a spectral scan and coversion to wavelength must be performed separately before
combining into a single spectrum. The dispersion is calculated for the
center of the array (row 128). Note: In the I band, operation at m=4
is significantly more efficient than at m=3 short of 1.06 microns and yields
higher resolution as well. However, order overlap restricts m=4 operation
to the range 0.90 <
The resolution of grating 2 is sufficiently low that the entire spectrum for
any blocking filter is contained on the array. Note: Order overlap
restricts I band operation to the range 0.90 <
Note: Due to the sharpness of the blace function for this grating, the
efficiency falls off significantly for grating settings > 1600 ECU. In
particular, the efficiency in the H band at m=3 falls rapidly beyond 1.65
microns, as the energy is directed into m=2, so use of this grating at H
is not recommended. Performance in the K band is also degraded beyond
2.35 microns, so the setting of 1595 ECU covers essentially all of the K
band within reasonable performance of the grating.
Note: Being blazed at 3 microns, grating 4 is significantly more
efficient in the I and H bands than grating 3. Moreovere, operation in
the I band at m=2 avoids order overlap with this relatively broad filter.
Interestingly, grating 4 is almost as efficient in the J band as grating 3,
with only sightly lower spectral resolution, so programs requiring both J and
H may be accomplished with grating 4, eliminating the need to change gratings.
Given the large variety of observing programs being carried out with the
Cryogenic Spectrometer, no single approach to data reduction is universally
applicable. This section is derived from a guide written by R. Elston and
reviews the promising approaches to data reduction for several specific
observing scenarios: observations of a bright stellar source, a faint stellar
source, and an extended object.
IR spectroscopy with the InSb array presents several problems not encountered
in normal long slit spectroscopy with an optical CCD. These include high
background with strong variable emission lines and
regions of bad pixels. To overcome the bad pixels, we recommend taking a large
number (>5) of observations of a given object and deriving the median of
the individual observations. True dark frames are generally
not necessary in the course of normal observing, since sky subtraction
removes the dark current as well, and flatfield exposures (except for L)
are best done by cycling the illumination lamps on and off. Most importantly,
the strong and variable background requires that sky subtraction be handled in
a two step process. First order sky subtraction is achieved by subtracting a
sky frame. Any residual sky due to variability can then be removed by
using normal longslit reduction techniques. The very high level of
IR sky brightness precludes the "traditional" approach of flatfielding
and background reduction on a single object frame, since even a 1%
uncertainty in the flatfield would produce effects large in comparison
to the object signal. The first order sky subtraction will typically
remove 95 - 99% of the sky, thus greatly reducing the effects of
flatfield uncertainties in the second order background removal.
The majority of experience with CRSP has been in the J, H, and
K bands, where the background consists of OH and O2 lines
superposed on the spectrum. Beyond 2.3 microns, the background
consists of a smooth continuum due to the emission from the telescope
as well as emission lines in telluric features. Thus, transitions which
produce telluric absorption lines in astronomical spectra will
produce emission lines in the sky background. Although this
scenario is more challenging, the same routines of sky subtraction
and off-object residual background subtraction
work well, as long as one operates in the linear (< 10000 ADU) region
of the detector.
Since the slit illuminates at most 138 pixels (cols 12 - 150), we recommend
that the raw images be trimmed to exclude nonilluminated portions of the array.
Not only does this yield smaller image files, but it avoids the undesirable
effects of dividing by near-zero values during the flatfielding stage. It is
critically important that all images be trimmed to exactly the same
subrasters to preserve pixel registration.
Linearization of the data is highly recommended, unless signal levels are
consistently < 5000 ADU. Data up to 13000 ADU (0.6v bias) can be effectively
linearized using the IRAF 'imred.irred.irlincor' task.
It is important that linearization be done on the raw
data, before any arithmetic operations on the data.
Finally, it is necessary to identify the dispersion as being along the column
direction in order for IRAF spectral extraction tasks to succeed. If the
entry DISPAXIS = 2 is not in the IRAF header, it may be inserted as follows
using the 'hedit' task:
It is necessary to produce flat fields using observations of the dome
White Spot, since the night sky possesses a complex spectrum in the IR.
To remove the effects of spectral features in the lamp and to achieve
relatively uniform illumination (and thus S/N) in the spectral
axis, we recommend that flats be made using the high dispersion
grating 1 for observations made with the low-resolution grating
2. Observations with gratings 1, 3, or 4 can, of course, be flattened with
the same grating, although the signal will fall to near zero at the
ends of the J band with grating 3. We suggest taking a large number (>10) of
flats in each spectral region and an equal number of observations with the
illumination lamps off, to remove dark current and stray light or thermal
background common to both sets of exposures. Some observers choose to move the
grating a small amount (10 ECU) between observations, although such observations
will require normalization before combination with the other exposures to
generate the flatfield frame. The illumination should be set sufficiently
low to yield 5000 - 7000 ADU in a 5 sec frame, to minimize nonlinearity
related effects. In the L band, the illumination lamps are not required,
and the dark slide should be used for dark subtraction. In general, a flat in
a given filter should be appropriate for observations within that bandpass,
although obtaining flats at the same grating settings used for observations
further minimizes sources of systematic error.
To make the flat, take the median (using 'imcomb') of the "lights off"
(or darks) and subtract it from each of the dome flats. Then take the
median ('imcomb') of the dome flats, while allowing the flats to
be scaled by the median of the frames. Somewhat better results may
be obtained by using averaging (with sigmaclipping) in the 'imcomb'
routine, since the straight median will force integer values. One
may try both methods and evaulate the differences, if any.
Finally, the median flat is
normalized. If the flat observations are taken at different grating
settings, use the 'response' task to remove low-order spectral slope
and to normalize the individual flats before median averaging to the final flat.
One may use 'imreplace' or other IRAF masking tasks to remove zeros or
negative values due to bad pixels from the flat. The final flat may be
normalized with 'response' if desired.
An important point covered in the manual is worth repeating. One must
be careful in the use of the multiple-read low noise code for the
flats (and bright standards), where the integration time is short.
The use of a read address time (315*N ms , where N is the number of
reads) which is a significant fraction of the integration time will
result in the collection of more charge on the array than indicated by
the output signal. This may inadvertantly push the operation of the
array into the increasingly nonlinear regime. A safe mode is to use
a single read for flats and bright standards. The higher read noise is
not significant, given the photon noise from the signal.
With the present gain of 7.2 e/ADU and read noise of 35 e, a signal level of
only 170 ADU will yield shot noise equal to the single-read noise.
Several approaches to wavelength
calibration are possible. One may use 'apall' to extract the sky spectrum
from the median sky frame, using the object spectra as reference objects.
One can then use the night sky lines to calibrate the wavelength scale,
using 'oned.identify'. This spectrum may then be assigned to the object
spectra with 'refspec' and the wavelength calibration carried out with 'dispcor';
refer to the documentation in the 'oned' package. Alternatively,
one may use the absorption lines in the object spectra as a wavelength
calibration grid. This has the advantage of being tied to the coordinates
of the object, which may not be the case for the sky spectrum, particularly
if the slit is wider than the seeing disk; in addition, sky lines may not
show up in short standard spectra. The disadvantage is that unlike the OH
lines, absorption features are usually a convolution of bands, with an
effective wavelength dependent on resolution. For grating 2, the dispersion/
pixel is essentially constant across the array, and it is usually possible to
set the wavelength scale by using 'splot' to locate a known
spectral feature in pixel units, calculating the wavelength of the first
pixel using the known dispersion, and setting the scale using 'p'.
For observations short of 2.5 microns, excellent results may be obtained using
the calibration lamps. This technique can have some advantages over the use
of sky lines, particularly for relatively bright sources, where the exposures
are often too short to yield good S/N in the sky lines, or in the short J
and I bands, as well as in the CO overtone region near 2.3 microns, where the OH
lines are very weak. The lamps provide good S/N in short observations, and
are particularly useful for removal of distortion (see the end of this
guide). The HeNeAr lamps built into the 2.1-m and 4-m guiders yield primarily
an Ar spectrum, with some He and Ne lines thrown in. Hardcopies of lamp
spectra with the lines identified are in the telescope domes, and the
"home$linelists/ar.dat" file in IRAF contains a list of these lines for
use with 'identify'.
For wavelengths beyond 2.5 microns, the thermal continuum from the calibration
lamps overwhelms the emission lines and saturates the detector very quickly.
For the L and M bands, one may observe the calibration lamp at the same
grating setting used for observing, but through the H or K filter, as
appropriate. It is then only necessary to multiply the resultant
wavelengths by a factor of two, using the 'hedit' task to modify the
header parameters "crval1", "cdelt1", and "cd1_1".
The best lamp spectra are obtained with the narrow slit 5, since this provides
better separation of closely spaced lines in the J and H bands and yields
a more narrow line for 'identify' to center on. With grating 2, many of the
lines will be severely blended and unsuitable for calibration. However, since
the dispersion is nearly constant across the array with this grating, it is
sufficient to pick out two or three unblended lamp lines and use a low order
fit in the 'identify' task. Although the use of slit 5 will result in a
small wavelength shift with respect to data taken with a wider slit, this
can be calibrated out as described in the next section.
It is necessary to correct spectra for telluric absorption lines by
observing a bright star near the source. This can be the same star
used for the suggested routine of determining the slit location
before moving to an object. A hot star will have a nearly featureless
spectrum with the exception of H absorption lines, so hydrogen line
programs may require modeling the continuum of the star
within the H line (c.f., Hall et al. 1981, Ap. J., 248, 898).
Typically, one should bracket program observations with
5 observations of the hot star and reduce as described below.
Normalizing the bright star spectrum and dividing into the reduced
object spectrum will produce a result normalized to a
Rayleigh-Jeans spectrum. For observations in the H band, where early dwarfs
display a large number of high-level Brackett lines, solar-type dwarf stars
may be preferable for telluric correction, although atomic lines of Si, Al, and
Mg will be seen with the resolution of grating 1.
A common problem in telluric removal is a small wavelength shift between
the object and standard spectra. This may result from flexure in the
spectrograph or (more likely) from the centroid of the object being not
centered in the slit; the latter is more likely when wider slits are
used. Since even a small shift in the wavelength calibration of the two
spectra will produce large effects in the vicinity of telluric lines, it
is important to precisely align the spectra before division.
In addition, the S/N must be sufficient to precisely identify a telluric
feature in all of the spectra.
In the real world, it is unusual to end up with a telluric or flux standard
at the same air mass for each object spectrum, so the result of the procedure
above may still contain residual telluric features. It is possible to generate
an "extinction spectrum" to remove most of these residual features using two
observations, well separated in airmass, of the same star. The two spectra
should be wavelength calibrated and coaligned in wavelength with the other
spectra in the group, as described above.
Using either 'sarith' or the interactive 'splot', generate the extinction
spectrum by taking the ratio of the two spectra, converting to log scale,
and dividing by the airmass difference in the spectra. Since slit losses will
make the absolute signal level unreliable, the result should be shifted
to give a value near zero where the atmosphere is known to be clean.
The correction process involves multiplying the extinction spectrum by the
difference in airmass between an object and its standard, converting to
dex, and multiplying this and the object/standard ratio.
This procedure has an associated element of trial and error. Because the
water vapor (and extinction) can vary over the night, one may have to try
adjusting the airmass difference multiplier to the extinction spectrum to get
best removal of the residual telluric lines. A more fundamental limitation
results from telluric lines which are saturated and thus do not obey a simple
extinction law. In practice, one may find that this technique can clean up
residual weak telluric lines, while giving poor correction for stronger ones.
Flux calibration of infrared spectra has historically been an informal matter.
The calibration tasks in IRAF utilize a database of optical spectrophotometric
standards calibrated in a non-logical manner which is not easily extendable
(should one even wish to do so) to the infrared. A simple approach which may be
utilized directly or upgraded as desired follows. Beginning with an object and
flux standard spectrum, generate the ratio as described above and correct
for the ratio of integration times. Taking the log of the result and multiplying
by -2.5 yields a spectrum of the object in magnitude units (alternatively,
one may have a photometric magnitude for the object from other sources, which
can be used to calibrate the magnitude spectrum). A simple IMFORT routine, such
as magfl.f, can convert the spectrum to
flux units, utilizing the wavelength information from the IRAF header. This
routine assumes a zero magnitude spectrum defined by a 9850 K blackbody.
To utilize a program such as this in IRAF, it is necessary to compile the
IMFORT version, yielding an executable 'magfl.e', and then defining this
to be a foreign task within IRAF by entering task $magfl = $magfl.e
(the '$' prefix indicates that the task does not have an associated parameter
file). Refer to IRAF documentation for more details.
One may define a bright star as one which can be observed in frame
times of less than 60 s; in this case, the greatest limitations to
S/N are systematic effects. To overcome these, take 5 - 7 observations
of the star, displacing it along the slit each time. The slitscan
script is ideal for this operation. Since the
observations span a short time, the sky will not vary significantly,
and one can generate a sky frame from a median of the individual
observations. Subtracting this median sky from the individual
observations removes the dark and first order sky. Each image is
then divided by the flat. At this point, one may
use 'apall' (in the 'twod.apextract' package) to extract 1-D
spectra from the flattened images. Second order sky subtraction
can be achieved by using the background subtraction in 'apall'.
The dispersion axis (column, or 2) must be set in the parameter set
for 'apextract'. The 'apall' task is apallingly complicated, but it is well
documented. To produce the final spectrum, take the median of the individual
1-D spectra, allowing scaling by the median within a sample interval where
the spectrum is reasonably constant.
A faint star is very similar to a bright star, except that the sky
will vary over the time of observation. One approach is to move the
compact source along the slit between observations, as with the
bright star, and use the observations before and after a given
frame as the sky. The average "sky" observation is subtracted from
the frame, and the result is divided by the flat. At this point,
use 'apall' with background
subtraction to remove residual sky and extract the spectrum.
As before, take the median of the individual 1-D spectra while
allowing 'imcomb' to scale by the median. Finally, divide the median
spectrum by the spectrum of the bright hot star to remove telluric absorption.
An extended object is one that fills much of the slit. In this case,
it is necessary to take bona fide sky observations between
observations of the object, although gutsy and/or impatient observers
will intersperse two object observations between sky frames. In
either case, it is a good idea to take several spectra of the object,
shifting it along the slit a small amount each time to eliminate the
bad pixels. To achieve first order sky subtraction, take the average
of the sky observations before and after and subtract from the object
frame. The result is then divided by the flat. Second order sky
subtraction can be achieved using the 'background' task in the
'longslit' package. When running this task, set axis=1 and
interactively select windows on either side of the source. The
task then fits a function between the windows and subtracts the
fitted value. The final step is to shift the
spectra along the slit to a common center. Use the 'identify' task
to find the peak of the spectrum ('identify' section = "mc" coord = "").
Placing the cursor on the peak of the spectrum
and typing m will give the centroid of the spectrum. One can
then use 'imshift' to shift the spectra to a common pixel.
Finally, use 'imcomb' to take the median of the shifted image.
If the source is reasonably bright, allow 'imcomb' to scale by
the median, while setting the sample section to the peak few
pixels of the spectrum. To correct for telluric absorption, divide the
2-D median spectrum by the 1-D spectrum of the hot star.
Sky or lamp calibration lines are noticeably curved. This is
a normal consequence of the grating equation in three dimensions, and is an
effect which is a function of the grating angle, being essentially
nonexistent with grating 2, weakly present with gratings 3 and 4, and significant
with grating 1, especially at the largest grating angles in the long
wavelength portion of the K band. This can result in a shift of as much
as 3 pixels between spectra at the middle and ends of the slit, an effect
which can degrade the resolution if several spectra taken along the slit are
combined. One may minimize these effects by either avoiding the end regions
of the slit or by separately calibrating spectra taken at different positions
on the slit before combining them. Where such steps are scientifically
impossible (e.g., a spectral image of an extended nebulosity which fills
the slit), it is possible to remove the distortion from the images.
The interested reader is referred to the manual "Reduction of Longslit
Spectroscopic Data Using IRAF" by Ed Anderson. The basic procedure is to
use a calibration spectrum (a sky spectrum with lines or a lamp) and fit
a wavelength solution using 'identify'. The task 'reidentify' will
repeat the identification process at different spatial positions in the
image, resulting in a grid of wavelength solutions as a function of
spatial position. The task 'fitcoords' will establish a "distortion
map" from this information, which can then be used to rectify images
with the task 'transform'. Examples of "before" and "after" are shown
in the Figure below. Note that the rectified image is also wavelength
calibrated.
Appendix IV: Suggested Grating Settings
grating 1 128
grating 2 177
grating 3 173
grating 4 169
Suggested Full-band Settings for Grating 1
Filter
Order
ECU
I
3
2000
0.8975
0.9599
1.0229
.000498
I
3
2220
1.0039
1.0634
1.1234
.000469
I
3
2350
1.0651
1.1231
1.1811
.000455
I
4
2610
0.8879
0.9285
0.9691
.000319
I
4
2760
0.9368
0.9760
1.0152
.000307
I
4
2920
0.9889
1.0261
1.0634
.000292
I
4
3080
1.0388
1.0743
1.1099
.000279
J
3
2430
1.1001
1.1565
1.2133
.000444
J
3
2580
1.1680
1.2220
1.2767
.000426
J
3
2760
1.2480
1.2994
1.3513
.000404
H
2
2080
1.4005
1.4924
1.5851
.000726
H
2
2260
1.5307
1.6181
1.7061
.000690
H
2
2430
1.6494
1.7335
1.8182
.000664
K
2
2880
1.9496
2.0244
2.1000
.000591
K
2
3070
2.0709
2.1406
2.2108
.000550
K
2
3260
2.1865
2.2515
2.3171
.000512
K
2
3450
2.2954
2.3569
2.1488
.000483
K
2
3640
2.4031
2.4590
2.5153
.000440
L
1
2100
2.8378
3.0277
3.2063
.001444
L
1
2340
3.1814
3.3622
3.5309
.001367
L
1
2580
3.5129
3.6840
3.8422
.001280
L
1
2820
3.8322
3.9939
4.1412
.001204
L
1
3300
4.4311
4.5295
4.6933
.001025
L
1
3500
4.6627
4.7913
4.9048
.000946
L
1
3700
4.8838
5.0022
5.1051
.000865
L
1
3860
5.0518
5.1619
5.2561
.000795
Settings for Grating 2
Filter
Order
ECU
I
4
850
0.7939
1.2391
.001746
J
4
900
1.0010
1.4424
.001730
H
3
950
1.3347
1.9232
.002313
K
2
875
1.7642
2.6538
.003490
L
1
700
2.4098
4.2185
.00709
Suggested Settings for Grating 3
Filter
Order
ECU
J
3
1345
1.0782
1.2177
1.3585
.001099
K
2
1595
2.0170
2.2252
2.4241
.001591
L
1
1195
2.7905
3.2289
3.6515
.003367
L
1
1345
3.3534
3.7824
4.1944
.003288
Suggested Settings for Grating 4
Filter
Order
ECU
I
3
1440
0.8952
0.9994
1.1036
.000818
I
3
1580
1.0022
1.1040
1.2057
.000798
J
2
1200
1.0656
1.2262
1.3869
.001260
H
2
1490
1.3996
1.5538
1.7079
.001209
H
2
1600
1.5251
1.6774
1.8296
.001194
K
1
1120
1.9465
2.2703
2.5942
.002540
L
1
1540
2.9176
3.2222
3.5268
.002389
Appendix V: Data Reduction Guide
INTRODUCTION
TRIMMING AND LINEARIZATION
FLATS
WAVELENGTH CALIBRATION
TELLURIC ABSORPTION
EXTINCTION
FLUX CALIBRATION
BRIGHT STAR
FAINT STELLAR SOURCE
EXTENDED OBJECT
DISTORTION CORRECTIONS
