||Perpendicular to slit
The 40 l/mm grating pixel sizes are independent of wavelength setting.
The 150 l/mm grating pixel sizes are virtually independent of wavelength
Because of anamorphic magnification, the Echelle spatial pixel scales vary with grating angle. The 0.41" probably changes with grating angle too, but I don't have measurements of that.
||Perpendicular to slit
Setting Slit Position Angles (PA)in CGS4
When setting the slit PA in the CGS4 software, it is important to remember:
The array is displayed `upside down' on the movie screen
so that south is to the top, and north to the bottom.
You need to consider in which direction the nod, or slide, should be specified.
Although you are asked to enter a PA which is east of north in your config,
you normally have to enter a negative PA. This is discussed in a little
more detail below.
The best way to demonstrate how to specify the slit PA is by showing some
examples which is done below. In each example, the top and bottom of the
array are marked, as well as north, south, east and west. An arrow indicates
the direction of a positive slide (or nod). Remember the telescope
moves in the opposite direction to this when you slide along the slit,
so be specific when giving instructions to the telescope system specialist.
Please be aware that, in order to try and show things as the observer sees
them on the array when the slit is north-south, north is shown at the bottom
of each plot.
Example 1: PA = 0 degrees (0 degrees east of north).
In this example, a position angle of 0 is entered into the config. the
slit will set to a north-south direction, with the south to the top of
the array. When using a positive slide, the offset beam appears to the
south of the main beam (or towards the top of the array).
Example 2: PA = -45 degrees (135 degrees east of north).
Here, a PA of -45 degrees has been entered. The top of the array is
now SE, the bottom NW.
A positive slide moves the offset beam on the array towards the SE
(i.e., the telescope moves NW)
Example 3: PA = -90 degrees (90 degrees east of north).
A PA of -90 degrees has been entered. The top of the array is in the
east, the bottom is in the west. A positive slide will move the offset
beam to the east, i.e., the telescope moves to the west.
Example 4: PA = -135 degrees (45 degrees east of north).
A PA of -135 degrees is entered. The top of the array is now in the
NE, while the bottom is in the SW. A positive slide will move the offset
beam to the NE, i.e., the telescope moves to the SW.
Some further notes on slit PAs
The default slide_slit command in the quad_slide exec is +11.4 arcseconds.
This positions the offset beam 19 rows up from the main beam. To move in
the other direction (towards the bottom of the array), simply enter a negative
value for the slide. Remember that the slit is 90 arcseconds long You normally
do not want to specify a nod which places the offset beam off the array.
Although the config demands a position angle east of north, it is only
possible to enter such an angle to a maximum of about 10 degrees. To set
to an angle greater than this, you must specify a negative angle. For example,
if you wish to set to a PA of 50 degrees east of north, set to an angle
of -130 degrees (simply subtract 180 from your desired PA).
The pixel size is 0.61 arcseconds. To make sure you nod onto another row,
you need to specify a value in arcesonds which is a multiple of this number.
The system usually does a fairly good job of positioning the crosshead
on the slit in the offset position. To increase the accuracy of the
offset position beam, you can do one of two things (or both): i) Use a
shorter slide distance than the default +11.4 arcseconds, or ask the
TSS to peakup in both the main and offset beams and correct the
current slide angle.
The current r.m.s. setting accuracy of the slit PA is 0.15 degrees.
Always inform the telescope system specialist when
you change the PA. The peakup position will change with PA, and if the
TSS is unaware of a change, time will be lost in reaquiring your target.
NOTES ON USING THE ECHELLE GRATING IN CGS4.
When using the echelle grating there are more details that need to
be checked than with the low resolution gratings. These are summarised
in this note.
1. Pixel and Slit Sizes.
Due to anamorphic demagnification, CGS4's square pixels do not view
a square patch of sky. Although the solid angle seen by a pixel is
constant, the shape is a function of grating angle. The difference
from a square field of view is very small at the low angles of the
first order gratings, but is quite prominent at echelle angles. For
example, at the nominal blaze angle of 64 degrees pixels are more
elongated in the spatial direction than the dispersion direction by a
factor of ~1.6. With the long camera at angles close to optimum the
pixel size is 0.91 arcsec (spatially) X 0.40 arcsec (spectrally). The
nominally 1-pixel wide slit is in fact about 1.3 pixels wide. The
2-pixel wide slit is 2-pixels wide (ie. 2arcsec spectrally).
The change in shape of the pixels as a function of grating angle is
very rapid at large angles. As a consequence, when using the echelle you
should always measure the size of the pixels by peaking up in two rows.
Edit your sequence to put in the correct offset between the rows
being used when nodding (see also the comment on slit alignment below).
The slit is slightly curved, so atmospheric and arc lines are not quite
perfectly aligned on one column of the array. In addition the dispersion
axis is also slightly curved, so that the spectrum of a point source is
not perfectly along one row of the detector. Both of these effects are
present with the low resolution gratings, but they are more significant
with the echelle. The curvature of the dispersion direction means that
the dispersion has slight dependance on row on the array. This means that
if an arc/atmospheric/astronomical line is straight along a column of the
array in the middle of the array, lines at the edge of the array will not
be so well aligned. In the dispersion direction the curvature amounts to
approx 0.5pixels across the 256 pixels. In the spatial direction the misaligment
with the columns due to curvature is about 0.1-0.5 pixels close to the
centre of the array, depending on grating angle. For point sources nodded
along about 30 rows the effects are very small and it is still reasonable
to combine the two nodded spectra to cancel sky lines.
There is software available, for example in Figaro, which was designed
to remove these sorts of distortion for optical spectrographs. There is
information about using these in the CGS4 data reduction notes in
curvature in CGS4 spectra . Using these techniques a residual distortion
of less than 1% can be obtained. However, in conditions of good seeing,
CGS4 data are undersampled in the spatial dimension and this may mean that
the effects cannot be fully removed.
The curvature of the slits means that even if the postion angle is nominally
N-S on the sky there will be a small (0.3 - 1 arcsec) E-W offset when you
nod along the slit. Since this offset depends on the grating angle and
how far along the slit that you choose to nod, it means that you should
measure it for each echelle wavelength/order that you observe at. After
you have measured the offset between the two rows you need to update your sequence accordingly.
INFORMATION FOR CREATING CONFIGS
3. Optimum Orders
When using the Echelle it is normal to select ECHELLE_AUTO_ORDER as the
grating when defining a config. If this is selected then the CGS4 software
will automatically select the optimum order for your wavelength. The optimum
order gives a higher throughput than other orders. The optimum ordering
software has recently been improved and now works very well for all wavelenghts.
You may wish to use another order, e.g. to get lower or higher spectral
resolution or to increase the spectral coverage. To use a different order
than the optimum one selected automatically, choose ECHELLE as the grating
when defining your config and then explicitely enter the order that you
You can calculate the optimum order for any wavelength as follows. The
blaze angle of the echelle is 64.6 degrees. This corresponds to the product
of wavelength and order (n x lambda) being about 55.0 microns. (e.g., we
find that the transmission at 2.12um in 26th order (n x lambda = 55.1)
is 10X higher than in 25th order (n x lambda = 53.0) In general for any
wavelength, the order for which n x lambda is closest to 55.0 gives the
highest efficiency. However, at lower orders (longer wavelengths), it is
best to fudge this somewhat, as the efficiency drops off more rapidly at
higher angles than at lower angles. The CGS4 software uses a lookup up
table calculated according to the above, with appropriate adjustments at
the longer wavlengths, to select the order for the echelle.
4. Order sorting
Longward of 1.3um, CVFs serve as order-blockers for the echelle. Shortward
of 1.3um, narrow band filters (1.083um, 1.233um, 1.257um, 1.282um), allow
echelle observations at important wavelengths. Wavelengths shortwards
of 1.3um may only be observed through these narrow band filters . The
longest wavelength observable with the echelle is 5.7um. The CGS4 software
will automatically select an appropriate filter or set the CVF to the requested
wavelength. It will report an error if your chosen wavelength is outside
the range of either the narrow band filters or the CVFs. Note that a long
wavelength leak is suspected in the short wavelength CVF (roughly H-band),
and is under investigation. If you want to do Echelle spectroscopy in the
H-band please check with your support scientist for an update on this situation.
5. CVF gradients
There is a slight gradient in the transmission of the CVFs along the slit.
The CVF calibration has been set for the middle of the illuminated area,
row 134. If you want to use rows towards the edge of the illuminated region
or nod more than about 30 pixels you may wish to check the CVF calibration
for the rows you will be using. First of all define an astronomical config
for the echelle in the normal way and set to this config. Peakup a star
on the desired row or look at a lamp line and then run MOVIE. Now while
MOVIE is running go into the menu called DIRECT_MOTOR_CONTROL . This menu
allows you to define an intermediate configuration. The menu items diplayed
represent where the CGS4 motors are currently positioned. To check the
CVF calibration try changing the CVF wavelength by a very small amount
and check whether the signal on the MOVIE display increases or decreases.
Once you know what wavelength you want to set the CVF to calculate the
difference between the grating wavelength and the desired CVF wavelength.
You can then use UKIRT_PREP to save an astronomical config with this CVF
offset. Be very careful if you decide to make a change like this - it is
possible to get "lost" in order on the CVFs because at 1-2.5um the Echelle
orders are very close together. For example 2.2um in 25th order is at the
same grating angle as 2.11um in 26th order - so if you move the CVF by
as much as 0.09um you could be looking at the wrong wavelength and order.
6. CVF fringes
Particularly in the thermal IR a ripple is seen in echelle spectra which
is caused by fringing from the CVF. This ripple can be difficult to remove
if there are amplitude variations between your source and your star. If
you observe very strong ripples, try moving the CVF by a very small amount,
or try puttting your source slightly out of focus. I think the latter helps
because it makes the source and the background fill the slit to the same
degree. Also take oversampled flats in preference to the usual undersampled
7. Wavelength Calibration
Because of the narrow wavelength range covered by the array when the echelle
is used, there are wavelengths where it is impossible to find lamp lines
that fall on the array. In this case you may be able to find lamp lines
at different wavelengths that are present at higher or lower order at the
same echelle setting. I.e., for such a line of wavelength lambda there
is an n that gives the same n*lambda as your observing setup, For example,
you want to observe at 3.00um in 16th order, where there is no line, but
notice that there is an Argon line at 2.40um which in 20th order would
appear at the same echelle angle. Such lamp lines can be found by (1) setting
the echelle to the wavelength you wish to observe (ie in this example 3.00um)
with 16th order selected and (2) changing the arc CVF wavelength to the
wavelength of the calibration line. The arc section of a config allows
you to enter a different CVF wavelength for observing arcs than for observing
your source. This arc CVF wavelength will only be used when you take arcs.
At some other wavelengths you will see lamp lines that were not expected
- these are strong lines in a different order and wavelength being transmitted
through the wings of the CVF transmission profile.
At wavelengths beyond the K window, observable lamp lines are generally
few and far between. For calibration with the echelle, one often must use
the above technique for finding arc lines in different orders, telluric
absorption/emission lines (atlases are available in the control room at
HP, and at JAC), or observations of astronomical line sources.
The Echelle and Peakups etc
There allways seems to be confusion about the Echelle grating and
the need to do 2 row peakups etc. Hopefully, this text outlines the
There are two complications with the Echelle that make this more
complex than with the low-res gratings, firstly the curvature of the
slit, and secondly, the fact that the pixel size varies significantly
with echelle grating angle.
Here's my suggested proceedure, assuming that you're nodding along
Initially, set up your ORAC Science Program with the default offsets.
Slew to the target, and ask the TSS to do a 2 row peakup, using
the SAVEPEAK and SAVE_PA commands. You will need to tell them which
two rows you would like them to peak up on. I suggest 140 and 160, the
same as we use for the low res gratings. It will help the TSS if you
can tell them the approximate offset in arc-secs between these two
rows. To do this, look up the grating angle on one of the VAX terminal
screens once you have set to your config, then go to the table in the
"Pixel scales" section of this document and estimate the pixel scale
for the grating angle that you're at.
The SAVE_PA command the TSS uses on the TCS tells the TCS the
correction between the demanded slit angle in your config and the
actual angle on the sky of the point joining the position of the star
when placed on the slit at the two points intersecting the peakup
rows. This means that you can simply ask ORAC to "offset along the
slit" and it will know the exact angle to go to hit the slit on your
two peakup rows.
Note that if you change to a different slit angle, there should
be no need to repeat the two row peakup. Simply do a
normal single row peak up on row 140 - the correction to the angle
will be exactly the same as before, assuming that you're still
offseting the same distance and that you haven't changed the wavelength
or order settings of the grating. If these conditions (ie same
configuration, same size offset) do not apply, then you will have to
repeat the 2 row peakup in your new configuration.
It is advantageous to your signal-to-noise ratio if you have the
spectra centred on pixel rows on the array. Peakup also helps you do
this. After doing the 2 row peakup, the TSS will be able to give you
the exact size of the offset between your two peakup rows. You should
go into the offset iterator in your sequence and replace the default
offset size (usually 11.74) with this number. Again - there is no need
to enter an offset component perpendicualr to the slit unless you
really do want to offset perpendicular to, and hence off, the
slit. This offset size will apply to all observations in the same CGS4
config, no matter what the slit position angle.
Notes on Configuration parameters
GRATING: Select the grating to use. Only two gratings are in
CGS4 at any one time - if in doubt your support scientist or telescope
operator can advise which they are. When observing with the Echelle
you should normally select ECHELLE_AUTO_ORDER.
WAVELENGTH: Enter the central wavelength for your spectrum in
ORDER: Enter the order to use the grating in. If you have selected
ECHELLE_AUTO_ORDER the value here will be ignored.
The CGS4 sensitivity page contains
information about the appropriate choice of orders for the gratings as
a function of wavelength. Normally echelle_auto_order should be used
for echelle observations. However see the echelle notes for details of where this
is not appropriate. If in doubt your support scientist can advise you.
CVF_OFFSET_(ECH_ONLY): If necessary enter an offset in microns
for setting the CVF when you are using the Echelle. Normally the default
of 0.0 should be used. This value is ignored for the 40 and 150 l/mm gratings.
The precise calibration of the CVF for Echelle observations can depend
on which row you choose to observe, the order and the slit width. Sometimes
it is desirable to set the CVF to a slightly different wavelength from
the grating to correct for this. Your support scientist will help you check
the CVF setting for your wavelength and order. Enter a non-zero value for
a CVF offset with caution and only after checking flats taken with NO CVF
POLARISER: The polariser should be NONE for normal spectroscopy.
Select PRISM or GRID if you want to do spectropolarimetry.
You cannot do spectropolarimetry with the echelle and prism.
SLIT_WIDTH: Choose a 1-pixel, 2-pixel or 4-pixel wide slit.
There is 1 pixel per resolution element and so a
two pixel wide slit will give reduced, resolution, but can be useful in
poor conditions. The 4-pixel wide slit is not available with the Echelle.
POSITION_ANGLE: Enter the position angle of the slit on the sky
in degrees EAST of North. 0 deg is a N-S alignment. Due to a historcal
reason, CGS4 will not accept angles greater than about 10 degrees. To request
a slit position angle greater than this, you need to specify an angle 180
degrees away. E.g., if you want 45 degrees east of north, you need to specify
-135 degrees as the position angle.
OBJECT_SAMPLING: The sampling required x the number of pixels
to sample over.
There is 1 resolution element per pixel so the detector is stepped to
fully sample. To Nyquist sample your spectra a sampling of 2 should be
selected (ie 2 data points per resolution element). Many users prefer to
slightly over-sample and choose a sampling of 3. To help eliminate bad
pixels from your spectra you can also choose to carry out the sampling
over 2 pixels. This means that each data point in your final spectrum has
been observed with two neighbouring pixels. When the data reduction interleaves
the raw integrations data from a bad pixel is replaced by that from its
good neighbour. Most observers therefore chose sampling of 2x2 or 3x2 for
As an example, if you choose sampling of 2x2 the data in the final spectrum
is taken at the following detector positions :
resolution elements | | | | |
detector pos1: 1 2 3 4 5 ...256
detector pos2: 1 2 3 4 5
detector pos3: 1 2 3 4 5
detector pos4: 1 2 3 4 5
Except for at the begining and end, or where there is a bad pixel, each
data point in the final array is the average of two measurements. For 2x2
sampling the reduced spectrum has 514 data points on the x-axis.
DATA_ACQUISITION_MODE: Choose ND_STARE, STARE or CHOP.
For exposure times less than 1 STARE is a little more efficient but
NDSTARE may aso be used. Use ND_STARE for all other observations that do
not require chopping. ND_STARE uses a multiple non-destructive read algorithm
to give optimum read noise performance.
SINGLE_EXPOSURE__SECS: Enter the on-chip exposure time in seconds
to be used for observations of your source. See the optimum
exposure times for a guide to what value to enter here. The minimum
exposure time is 0.12secs for the 256x178 array. For the 256x48 array,
the minimum time is 0.023 secs, and for the 256x32 array, the minimum time
is 0.016 secs.
EXPS/INT_OR_EXPS/CHOPBEAM: Enter the number of exposures to be
coadded into an integration (integ) at each array sample position.
For long exposures on faint sources set object_exp_per_integ to 1. For
very short exposures on a bright source it is a good idea to choose a value
so that the time per integration is a few seconds for best observing efficiency.
CHOP_CYCLES/INTEGRATION This only applies to CHOP mode, and is
ignored otherwise. Enter the number of chop cycles you want before nodding
the telescope (one chop cycles is an object-sky pair).
FLAT_SAMPLING: choose AS_OBJECT in most cases
For Echelle users it is now possible for the data acquisition to take
over-sampled flats with the same sampling factor as your object observations
to help remove CVF ripple.
FLAT_LAMP: choose your required lamp. The numbers refer to black-body
apertures. There is a table of flats
of appropriate values as a function of grating wavelength and order.
If you choose off then the calibration unit will not be used - useful for
if you want to observe the sky, or the dome etc for a flat.
FLAT_ACQUISITION_MODE: Normally this will be NDSTARE
FLAT_EXPOSURE_SECS: typically 0.12 - 1 sec. See the
table of flat exposure
FLAT_EXP_PER_INTEG: The number of exposures to be coadded into
an integration. typically 50-100.
FLAT_INTEG_PER_OBS: The number of integrations to be taken and
returned to the Vax. Typically 1.
ARC_LAMP: Choose your required lamp. There is a table
of recommended arcs lamps as a function of wavelength. Plots of sample
arc spectra can also be found at that page. If you choose OFF for the lamp
then the calibration unit will not be used - useful for if you want to
observe sky lines for wavelength calibration.
ECHELLE_ARC_CVF_WAVELENGTH: Defaults to the same wavelength as
for your object observation.
Because the wavelength coverage on the array with the Echelle is not
very large it is possible to observe at wavelengths where there are no
lamp lines which would fall on the array. This allows you to select a different
CVF wavelength from the grating, to look at lamp lines in a different order
from your source.
If the value shown here is -1 the default CVF setting - ie the same
CVF wavelength as for your object will be used. If you have just recalled
or created a CONFIG then the default wavelength should appear here.
ARC_ACQUISITION_MODE: Usually NDSTARE
ARC_EXPOSURE_SECS: Typically 0.12 secs with the 40 l/mm grating
and long camera.
ARC_EXP_PER_INTEG The number of exposures to be coadded by the
into an integration at each detector sampling position. Typically 5.
All ARCS are automatically taken with the same sampling and sampling
range as your object observations. Although on-chip exposure times are
short, you should add sufficient exposures so that the time per sampling
position is at least 0.5 secs, to prevent the background increasing due
to the detector translation stage getting warmer.
ARC_INTEG_PER_OBS The number of integrations to be taken and
returned to the Vax at each detector sampling position. typically 1.
It has come to our knowledge that the 256 x 256 InSb array in
CGS4 suffers from small scale latent images. This problem has only
become noticeable since long integrations have become possible due to the
small pixel sizes in use with the long focal length camera.
In general, the latency only affects long exposures of very faint sources
after previously observing a bright source (e.g., a standard star). The
level of latency varies, but is usually less than about 5 counts (30 electrons)
above the general background level.
The simplest method to rid the array of the latent image is to ask
your TSS to run manpeak to take a few short exposures without changing
the apperture, once you are pointing at your faint source.
If your targets are bright enough to peakup on, then the above procedure
is unnecessary. The process of peaking up will remove any latent image.
CGS4 Object acquisition Procedures
Peak-up procedures with CGS4 seem to be a large source of
confusion, especially for observers determining before hand what their
requirements are in terms of astrometry, co-ordinates and reference
stars. This page aims to clear things up a bit.
The complication arises in that there are many ways to attempt to
ensure that you have accurately hit your target with the slit, all of
which apply to different situations.
Concepts and Introduction
The aperture defines the position of the slit within the
telescope focal plane. The default aperture is re-determined during
engineering time whenever the instrument has been off the telescope,
and at regular intervals inbetween. The aperture varies slightly with
flexure etc. Fine adjustments to the aperture definition are made
with the peak-up procedure.
The actual peak-up procedure is the process whereby the
aperture definition is finely adjusted to give the maximum signal on
the CGS4 array, whilst pointing at some (bright) target. This
involves taking several (typically 15 or so) CGS4 integrations. Each
of these must convincingly detect the object being peaked up on. It is
possible to automatically subtract sky frames from the images used
for peak-up. The integration time needed depends on the instrument
configuration in use and on how bright your target is.
With the 40l/mm grating, peaking up on a target fainter than 14th
magnitude takes a long time. If you need to estimate peak-up
times for fainter targets or other configurations, refer to the CGS4
sensitivity page, and use the following: Calculate the exposure time
for a 3 sigma detection of a point source at your target's
magnitude. Multiply this by a factor 2 (to account for the fact that
for most of the frames, you're not peaked up, so you're only seeing a
fraction of the light). Multiply this by 15 (say 15 observations
The Fast Guider is mounted on the cross-head. The
cross-head can position the fast guider very accurately (~0.1") within
the telescope focal plane. The crosshead limit is roughly a circle,
180" from the pointing centre (ie centre of the slit).
The Fast Guider brightness thresholds also vary with
conditions. Because the fast guider works in the optical is is
affected by scattered moon-light. In exceptional conditions -
clear skies, exceptionally good seeing and no moon, we can guide on
stars down to V magnitude of 18.7. A V mag limit of 16 or 17 is more
realistic for average conditions. On fainter targets, we increase the
integration time of the autoguider CCD, which reduces the Tip-Tilt
frequency. On bright targets, we guide at 100Hz. We can go down to
~20Hz for faint sources. Bright sources are more likely to give you
really good Tip-Tilt performance.
UKIRT employs a dichroic tertiary mirror, feeding the IR light to the
instrumentation, and the optical light to the fast guider. This means
that it is possible for the fast-guider and the IR instrument (eg
CGS4) to observe the same target.
Trivial case - bright (~14th Mag or brighter in the IR) guidable targets.
In this case you guide on the target, and you peak up on it before you
start taking data. The peak-up takes a few minutes.
No special co-ordinate, astrometry or guide star requirements.
Bright (14th Mag or brighter in the IR) non-guidable targets.
Reasons they might be non-guidable include: Target not point-like in
the optical, or target is so red that despite being bright in the IR,
it's too faint for the fast guider to see (see notes above).
In this case, you guide on an offset guide star (needs to be guidable
and within 180" of target), and you peak up on the target.
Requirements: A guide star. You can pick one using the ORAC-OT
preparation tool when you get to Hilo. You don't need any special
astrometry. If your co-ordinates are off by more than an arc-sec or
two, peak-up will take longer.
Intermediate brightness targets.
OK, the definition of "intermediate brightness" is non trivial. By
this, I mean something you can guide on, but something that
would take too long to peak up on. See the notes above about the fast
guider sensitivity and peak-up times.
Procedure: The TSS will slew the telescope to a nearby (within a few
degrees) bright CMC star, and will peak up on that whilst guiding on
it. This ensures that the slit and fast guider are at the same place
within the telescope focal plane. Then they'll slew back to your
target and guide on it. Because peaking up on the CMC star aligned
the guider and the slit, the slit will now be accurately placed on the
Note: The guider and slit will not move relative to each other over
a small slew to the CMC and back. If it were a long slew, there would
be concerns about flexure and differential refraction between the
optical and IR beams.
Requirements: The TSS will locate a suitable CMC star as needed. You
need reasonably accurate co-ordinates (ie within a few arc-sec). If the
target is towards the fainter edge of the guider's capability, better
co-ordinates will help. If the source is not point-like enough to guide
on, treat this as a faint target.
By "Faint", I mean too faint to guide on. This obviously implies also
too faint to peak up on. There are several options in this case.
This is the most reliable way of acquiring very faint targets. You
need an accurate astrometric offset between your target and a
nearby (within 180") guide star (see notes above on fast guider for
what qualifies as a guide star). You would usually measure such
offsets from a deep image of your target field. By "accurate", I mean
to within a fraction of the slit width you will be using. Say
0.2" for the 0.6" slit. The more accurate your offset, the more light
you'll get down the slit.
Proceedure: Simply enter the RA and Dec co-ordinates for the target
and guide star in the OT. Ensure that the co-ordinates you enter have
the correct offset between them, taking into account the factor 15
between seconds of RA time and seconds of arc. When you slew to the
field, inform the TSS that you have an accurate crosshead offset set,
and thus to adjust the telescope position rather than the crosshead
position to bring the guide star into the fast guider if necessary.
The TSS will ensure that the telescope pointing model is good prior to
you starting to take data. This will involve at least going to a
nearby CMC star, and probably doing a peak-up on it. If necessary they
will collimate the telescope pointing model to the CMC star
co-ordinates at this point.
Requirements: Accurate astrometric offset to a nearby guide star.
This is the "last resort" method, though the UKIRT Telescope Control
System and Pointing model is sufficiently good that this method
usually works well.
You would use this method if you don't have the astrometric offset to
a guide star necessary to carry out the previous method. You do,
however, need accurate absolute co-ordinates. You shouldn't rely on
this method with the 0.6" slit. Many observers have had great success
using this method to observe radio sources, with accurate co-ordinates
from high resolution radio interferometers, using the 1.2" slit.
Proceedure: Type your accurate co-ordinates into the OT. Use the OT to
find a nearby offset guide star. The TSS will go to a CMC star close
to your field and collimate the telescope pointing model. When they
slew back to your field, inform the TSS that the target co-ordinates
are accurate, and to move the crosshead to bring the guide star into
the guider if necessary.
Threats: You need accurate, co-ordinates. You might need to be
aware of the co-ordinate frame your co-ordinates are referenced to
(how well is this frame tied to the Hipparcos or Radio reference
frames?). This method becomes less reliable in high wind conditions -
during the time the TSS is adjusting the guider position to acquire the
guide star, the telescope is sat on it's main drives, as opposed to
being locked onto a guide star. If wind gusts are knocking the
telescope around at this point, the drives are unlikely to be able to
hold it in position to great accuracy.
Requirements: Accurate absolute co-ordinates.
If your targets are too faint to guide on, and you don't have
either sufficiently accurate co-ordinates or astrometric offsets, then
your best option is probably to use some time at the start of your
observing run to get UFTI images of your fields, from which you can
measure offsets in order to do a blind offset. You should discuss this
with your support scientist and Paul Hirst, the CGS4 instrument
scientist, well in advance of your arrival in Hilo.