Each telescope has a telescope control computer (TCS). On the 4-m and 1.5-m this is
a VME bus computer with a program written in C. The 0.9-m has a DFM Engineering
TCS. The TCS computer will be operated by your Observing Assistant at the 1.5-m
but at the 0.9-m you will probably be on your own. Refer to
"The 0.9m
Telescope"
All hardware functions in the CCD camera head (with the exception of the CCD shutter
are controlled by a STD bus computer (actually a PC) mounted on the
telescope. The operation of this computer should be transparent to the user.
At the 4-m and 0.9-m, the light is reflected from only the primary and
secondary mirrors before passing through two filter bolts, and finally a fused silica window
before reaching the CCD. On the 1.5-m f/7.5 and f/13.5 there is a fused silica corrector element
above the filter bolts. Each filter bolt can hold up to seven 3 inch square filters, which
must be less than 1 cm thick, plus a clear position. There are 3 inch square UBVRI sets
for use with the Tektronix CCDs. A single 5-position module which holds up to five 4x4
filters can be used if need be. In all cases the CCD is focused by moving the telescope
secondary mirror. The filter bolts are contained in a shutter-filter unit, which also holds
dark slides (manual), preflash arm and light-emitting-diodes (leds), plus a shutter.
Your telescope will be operated by an Observing Assistant who is familiar with the
operation of both the telescope and the CCD system. The Observer Support section will
have helped in the setup and check out the instrument, and they will respond to any
problems that might arise. The telescope should point to better than 10 arc sec rms (0.9-
m, 15 arc sec rms) assuming accurate input co-ordinates. If the pointing is much worse
than this then you should complain; on the 0.9-m and 1.5-m the collimation co-efficients
may not have been re-entered following reset of the Telescope Control System. It is a
short (10 min.) procedure to re-determine the collimation coefficients).
Peltier cooled CCDTV cameras are used for acquisition and guiding. The configuration
is different for the three telescopes. At the 0.9-m the TV can be scanned in 1-D over an
area approx 12x3 arcmin. The useable field is a little less than 4x3 arc min. At the 1.5-m
the TV is mounted on an X-Y stage ("the GAM") which can be moved until a guide star
is found. At the 4-m, the TV views a fixed position of sky via an f/2.5 focal reducer. The
guiding is done in a PC via a program written by Steve Shectman (LCO).
Observers needing further information should consult the following persons, in the first
instance contacting people who are actually on Tololo (Tol) at the time:
The IRAF-based user interface allows observing commands to be sent to Arcon from
within the IRAF cl. This results in a single uniform user interface for data taking and data
reduction and allows the Arcon user to employ features of the cl such as the parameter
mechanism and the history editor. It also allows advanced users to write cl scripts which
freely mix data acquisition and data reduction operations. We have tried very hard to
make the user interface look very similar to "ICE" as used at KPNO, and it should
become more similar in the future. The underlying hardware is, of course, very different.
So users should be aware that it is NOT identical and parameter lists, etc. will differ.
Below we describe recommended observing procedures. You are recommended to use
IRAF for reducing your data, and for writing your frames to dat or exabyte.
Information on the use of CCDRED, and IRAF in general, is given in the manuals and
cookbooks kept near the SUN terminal. Observer Support personnel and the Data
Reductions Specialists can assist you with IRAF. But you should not go to the telescope
and expect to learn IRAF at the same time as you observe!
THE MANUAL
"THE
ARCON-IRAF INTERFACE - A PRELIMINARY USER'S GUIDE FOR DIRECT
IMAGING"
Describes observing software and procedures. You should refer to this manual
for details on how to take data. Only a brief
outline is given below.
The Arcon controller boot-up is automatic (watch the Arcon console window). See the
Observing Software Manual for procedures. IRAF should automatically start in the two
IRAF windows, load the arcon package and connect itself to the
controller.
As with ICE, most of the details are hidden in four parameter files, obspars,
instrpars, detpars, and telpars. You will need to set these up at
the start of your run, afterwards they should need little or no modification. In particular
(and if you need to re-load the software, you should run setdetector force=yes
in order to synch the controller to the parameters you have set in detpars.
3.2.2 Data Storage.
You should check you know where your home directory is (show home) and
where your pixel files are going (show imdir). Check that you have lots of disk
space (disks). Make sure that when you do an observe that an IRAF
image ends up where you expect it do be. Display the image in IRAF, using
display. See below for instructions on how to check whether your CCD system
is functioning correctly.
3.2.3 CCDs with Multiple Outputs.
Some of the CCDs now converted to Arcon operation have more than one (usually
four) output amplifiers available. The CCD is read out through each amplifier in parallel,
the data is then digitized through (1,2,4) 16-bit analog-to-digital converters and then
transmitted on a single serial fiber-optic line to the SUN computer to a spool file, and
then converted to an IRAF image file. A raw quad-amplifier picture looks a little unusual.
Each quadrant has a slightly different electrical offset (ie overscan level), and the four
overscan segments are in the center of the picture. A CCD read out through two
amplifiers can either have two overscan strips at the right hand side of the picture
(parallel split) or two parallel strips in the center of the picture (serial split, eg Tek 2048
#3)). The real time display automatically removes the overscan and applies an offset to
each quadrant to normalize the DC level. In addition, the real time display colors any
saturated (ie 65535) pixel in red.
3.2.4 Processing Quad Images.
"Quad" CCD pictures are trimmed and overscan-subtracted in a separate pre-
processing task called quadproc (in the quad package). You will need
to load this package into your IRAF DATA REDUCTION window. Run
setinstrument to set up quadproc. Note that even after the pictures are
trimmed and overscan subtracted the sky levels in each quadrant will not be identical
(unless sky level is zero). This is because each of the four CCD amplifiers has a slightly
different gain. The procedure of quadproc'ing the pictures, plus the usual zero (bias)
subtract, flat-fielding etc is now all been integrated into quadproc.
3.2.5 Writing Tapes.
There should be an exabyte drive and a DAT drive (at the 0.9-m, ctio36!mtx)
ctio36!mtd) available in your dome. The Arcon produces
16 bit data of type 'unsigned short integer' with valid data ranging from 0 to 65535. Thus
the unprocessed frame from a 2K CCD requires about 8.8 Mbytes.
ccdinfo
and you change the gain by doing setdetector force=yes. For direct imaging gain
= 1 (2 for Tek 2048 #3) is in almost all cases the best to choose. The gain and read noise
are calculated using the gfind task which is in the contributed package.
You will first need to do two low level (bias) exposures, and two high level exposures
such as dome flats (both at the same light level).
You can also choose to bin the CCD if you wish. Given that the read-out times are short,
even for the Tek 2048, this is not a very popular option in general. If you wish to read out
a region-of-interest rather than the whole CCD, this is also an option, and you can
change from reading a quad amplifier CCD to using only 1 or 2 amplifiers if need be. All
these options are controlled by parameters in detpars. You should consult with
Observer Support or your Support Astronomer if you wish to make changes, not all
options are possible for all CCDs.
It is your responsibility to check that the instrument is operating to your
satisfaction.
A few tests are suggested below which should aid you in this evaluation.
You should also be on the lookout for any changes in the CCD characteristics during the
run, which may signal the onset of CCD or dewar problems. Count rates from dome flat
field (white spot) exposures should be the same from day to day, to better than 10
percent. You should also divide the night's flats by the previous night's and check to see
that the division is quite flat, with perhaps only the addition of one or two "rings" due to
specks of dust falling on the filters or CCD window (ask for these to be cleaned if the
number of spots seems excessive, but don't try to use your old flats later on). Examine
bias exposures, they should not show horizontal or diagonally slanting noise "stripes",
which move position from frame to frame. Sometimes very low level noise features are
seen on bias frames, before complaining check whether these are at a significant level
to affect your data. Most of our CCDs have very low level fixed pattern noise (1 adu or
less). Use of the IRAF Imhist command is also recommended. Histograms
should exclude bad columns, edges etc., they should be smooth with no missing codes.
It is worthwhile making a few test exposures before proceeding further. These tests can
be repeated as desired throughout your run. All can be done during the daytime.
3.3.2 Zero (Bias) exposures.
These zero illumination exposures should show no
noise patterns down to about the 1 adu level. They should be approx. flat although bright
columns will show up (Tek 2048 #3, STIS 2048). Now do an exposure (dome flat) that
adds a small signal. Any exposure with some illumination (preflash, flats) should show
the overscan strip on the right hand side of the picture (single channel readout, or down
the center (quad readout). Don't forget that the Real Time Display automatically clips out
the overscan so you will need to display the image within IRAF.
3.3.3 White Spot exposures.
Ask for the telescope to be pointed at the white spot
on the dome, and turn on the quartz halogen lamps. Make sure the telescope focus is
approximately correct for the filters you will be using during the night. Put the color
balance filter in the beam. This filter makes the 3000 C quartz lamps look more like the
5500 C night sky. Take a trial exposure, typically a few seconds through a broadband
(B,V,R,I) filter. Display the frame, and do a row plot across it. The response should be
reasonably flat (say 10 percent in general). After (quadproc if you are using a quad or
dual readout CCD) subtracting off the bias exposure you can try dividing frames at two
different count levels whereupon all the structure seen in the dome exposures should
disappear to better than a few 0.1%.
3.3.4 Dark Exposure.
A five minute dark exposure (lights off in dome) with no
preflash should only show the hot columns and hot pixels and a few cosmic rays. The
general background should be at the same level as the overscan. Cosmic rays will be
slightly trailed column-wise with a TI CCD otherwise most events should not show trails.
For all these CCDs, the dark rate is so low in comparison to sky background rates as not
to warrant doing long dark exposures.
3.3.5 Shutter Timing.
The shutter unit has a 100mm clear aperture. The effective
"open" time is always longer than the nominal number of seconds requested for your
exposure time. This is a delay prior to opening and closing. The speed of the shutter
blades is high and the corner-center exposure difference is only a few ms for the smaller
CCDs. For the Tek 2048 the center-to-corner difference is some 50 ms (ie exposures
are 50 ms longer than nominal at the center, and 0 ms longer at the corners). So beware
when doing short exposures for flat fields and standards. Sometime during your run, it
is a good idea to determine this correction. The easiest way of doing this is to compare
(for example) a 20 second dome flat with the sum of 20 1s dome flats. The latter can be
done extremely quickly using the focus procedure but not bothering to move the
telescope, whereupon you can avoid reading the CCD out 20 times but the shutter will
have opened and closed 20 times. Bracket this exposure with normal 20s exposures (just
choose a filter and exposure time to give several thousand adu total counts), the
bracketed exposures can be averaged to take out any drifts in the lamp output, but in any
case it is wise not to begin this test until several minutes have elapsed after turning on
the lamps. Subtract bias off the two frames, then if the counts are C1 for the sum of the
20 1s exposures, and C2 for the single 20s exposure, then to first order we have
shutter delay time (s) = (C1/C2) - 1.
You can easily do a contour plot which is probably the best way of displaying the result.
You could also in principle correct all your data. Remember -- you are likely to
compromise your photometry with the Tek 2048 if you take flat or object exposures
shorter than 5 seconds and do NOT take account of the radial variation of the shutter
timing.
Shutter times less than one second are possible, but below about 0.1s shutter operation
is unreliable. Non integral shutter times are possible, and are correctly recorded in the
image header. They will most likely be seen when using pause and
resume.
3.4.2 Zero (bias) Exposures.
These zero exposure can be started using the
commands zero or observe. You should obtain between 5 and 25
exposures, if you are doing really low S/N photometry through narrow band filters then
you should aim for the higher number. These can be combined into a single frame with
an "average sigma clipping" or "ccdclip" algorithm (IRAF), the idea being to produce a
"zero level" frame from which any cosmic ray events have been removed. You will get
plenty of cosmic rays on the Tek 2048 biases due to the relatively long readout time and
large area, so make sure you have at least 5 biases for this CCD
3.4.3 Dome (white spot) Flat Fields.
These exposures map the sensitivity variations
of the CCD, and since the response alters as a function of wavelength it is necessary to
prepare a flat field for each filter. Ultimately, the quality of your data will depend on how
well you prepare your flat fields. You can use the dflat or observe
commands to take dome flats. Broad band (eg BVRI) dome flats require use of a color
balance filter. Do not use the color balance filter for U, or for medium band (eg
Stromgren, DDO) or narrow band filters. You should aim for flat levels of several
thousand (say 5000 - 40000) counts. Take a sequence of exposures per filter
The resulting frames can then be averaged to produce a single dome flat field frame for
each filter. Data frames should successfully be flattened with dome flats to 1-2 %,
perhaps a little better for redder passbands such as V,R,I. To do better than this there
are two alternatives -either to use twilight flats alone, or to correct the low spatial
frequencies in the dome flats by using a dark sky exposure(s) or with the twilight flats.
However you are advised to take dome flats anyway, just in case the twilight or dark sky
flats are unsuccessful. You may find the dome flats are good enough for your purposes.
NOTE - You can get excellent pixel to pixel statistics with dome flats, and in some cases
this may be more important than low spatial frequency errors
3.4.4 Twilight Sky Flat Fields.
Broadband B and U (especially) twilight sky flat fields
appear to flatten data frames better than do dome flats. Unfortunately, twilight is short
and the number of flats you can take is limited, which can limit S/N. CCDs up to
1024x1024 can be read in just a few seconds, but with 2048's, particularly if you are using
several filters, other techniques include using smoothed twilight sky flats to take the low
spatial frequencies out of your dome flats, thus preserving the good high spatial frequency
S/N of the dome flats, or consider taking 2x2 or 4x4 binned frames and later use them as
smoothed corrections to your dome flats in order to correct the low spatial frequency
errors in the latter
You will need to point the telescope towards the east (HA -3 hrs) in order to lessen
polarization effects (in theory -- or just leave it at the zenith, polarization effects seem
to be minimal). Do not use color balance filters. As the sky darkens after sunset, take
one second exposures (with U if you are using it) until the sky is dark enough so that an
exposure of a few seconds gives about 20000 counts. Use the preview command,
and look at the counts in the Arcon STATUS window. You should have time to take
several flats in each filter, and if so these can be combined later (IRAF combine
in the ccdred package). To make it easier to remove star images from your flats,
take the flats with the TELESCOPE DRIVE ON, and jog the telescope a little after each
exposure. If you wait until it's too dark, then your exposures will be longer and the
chances of getting significant contamination by stars is greater, particularly in R and I.
Twilight flats should be done in the order UBVRI. Exposures should be at least 5
seconds long in order to reduce the center-corner variation due to the shutter to
negligible (<0.1%) proportions. Start the test U exposures at sunset or just a few minutes
after.
3.4.5 Dark Sky Flats.
Exposures of "star-less" high galactic latitude fields can be
used to provide an "illumination correction" to flat fields. The technique is to take several
stepped (e.g. 20 arc sec) exposures of the field then combine the exposures in such a
way (e.g. median filtering) that faint stars are removed. Since the final frame will be
heavily smoothed a sky level of 200-300 ADU over bias is sufficient. Be sure to remove
any bad pixels before smoothing. The resulting smoothed frame can then be combined
with the corresponding dome flat to form a master flat. That is, the sky flat removes low
spatial frequency errors in the dome flat. The advantage of this procedure over using
twilight sky flats is that you retain the good pixel to pixel statistics of the dome flats, and
the color match is improved. In practice the latter advantage does not seem to be very
important. The disadvantage of using a combination of dome flat + smoothed sky flat is
that any high spatial frequency errors in the dome flats will not be removed, however in
general these are few
Note that dark sky flats are necessary to remove any fringes due to night sky emission
lines causing interference in the thinned CCD (Teks). None of these CCDs fringe
excessively (<1% in the I band is typical) and so you may not need to de-fringe
It is also possible to prepare flat fields from combinations of your object frames. This is
generally only possible if they contain few bright stars, and even then it may be necessary
to "patch" out the brighter objects prior to combining the frames. If you are unfamiliar with
these rather specialized techniques you are encouraged to consult with CTIO staff prior
to your observing run.
3.5.2 Focusing the Telescope.
On the first night set the telescope focus to the
"nominal" value. Ask your Observing assistant or look amongst old Observing Record
Sheets to ascertain the relevant value. If the focus is far from correct, read the CCD out
using movie while you adjust the focus. Choose a focus star with V 10-11 (0.9-
m, scale to other apertures) whereupon your exposure time will be a few seconds.
Accurate focus is found by taking several snapshots of the star, each at a different focus
position. Use the focus command, with the relevant obspars parameters
set. It is normal to choose shtype = detector (charge will be reverse-clocked
rather than the telescope moved between exposures), and at the 0.9-m and 1.5-m
focmode=manual, ie you need to change the focus manually between exposures.
NOTE in all cases the double-spacing is made between the first and second exposures.
And for shtype = detector, on quad readout CCDs the stars in upper and lower
halves of the CCD will move in opposite directions.
Please save your focus frames and write a note on the night report sheet saying where
they are, they will be written to tape and kept as a database for seeing and optics
improvements. Offsets for the standard UBVRI sets are known (ask Observer Support)
so you need only focus one filter.
When focusing, always focus against gravity: move out (to larger readings of the
pedestal). The procedure is as follows:
Once rough focus is known, move focus in (lower numbers) about
500 units (0.9-m) to take up backlash. Exact focus is determined by doing 8 - 10
exposures at different focus settings, offsetting the telescope slightly between each one.
The CCD will be read out after the last exposure. Vary the focus by 40 units (0.9-m,
Futaba) between each exposure.
The quickest way to examine your images is to use imexam to plot profiles. Best
focus will be that with highest peak intensity and minimum distortion in the wings. This
last criterion is perhaps the most stringent one. The focus will change during the night,
as it appears to be a function of temperature and possibly telescope position. In good
seeing (FWHM <1.5") frequent focus checks are recommended especially if you move the
telescope more than 10 or 20 degrees.
3.5.3 TV Camera.
A Peltier-cooled CCDTV camera serves both for field acquisition and for guiding. Field
of view is approximately 4 by 3 arc minutes. On the CCDTV control panel integration time
can be switch selected and gain/offset controls are available. Try integration time of 1-2
seconds, gain = 2 and offset so that you see star images.
Focus the TV after focusing the CCD. Watch for smallest images, preferably with the
Leaky Guider switched to "digital". This focus is not that critical, as the guider will work
fine on slightly out of focus images.
Normally the determination of the TV-CCD offset will be carried out entirely by Observer
Support personnel on the first night only. In any case, in these days of large CCDs it is
more efficient to find objects and center up using the CCD in preview mode. But
if you want to determine the offset, see the Software Manual.
3.5.4 Guider Operation.
We are changing over to using a program written by Steve
Shectman, which runs in a PC and displays the star-field on the VGA monitor. At the
1.5-m, the TV is feed by a moveable guide probe. Be careful not to get the probe
vignetting the beam. For the Tek 2048, allowable guide probe co-ordinates are X =
anything, Y > 50.
3.6.2 Exposure times:
All the CCDs bleed charge down columns, and to a lesser
extent across rows, when an integration is long enough for a star image to become
saturated. All recover well from over-exposure, however excessively bright stars should
be avoided. All the CCDs have very low readout noise so there is an advantage in
dynamic range by splitting a long exposure up into a number of shorter ones. You should
also beware of bright stars just outside the CCD field. These can easily scatter light onto
the CCD frame and make photometry difficult.
Cosmic rays tend to be either single or double hot pixel events. The standard way to
remove them is to do at least 3 exposures (4 or more give better statistical results) of the
same field and median (or use avsigclip option in the IRAF combine command) filter to
remove transient events. The IRAF cosmicray and imedit commands can be effectively
used to remove cosmic rays from single frames.
You should be careful not to over-exposure your object(s) of interest. Saturated pixels
are colored RED on the real-time display. Beware that for 2048 CCDs the display
averages four CCD pixels for each display pixel. It is not (yet) clever enough to color the
display pixel red if one or more of the component CCD pixels is saturated.
3.6.3 Bad pixels:
there are two types: 1) hot (above local background) and 2) cold
(less than local background). Sometimes hot pixels tend to be found at the beginning of
a column of cold pixels - it appears that most of the charge finds its way into just one
pixel. Neither type are transient events: they are both permanent defects in the CCD.
However, the number of bad pixels is a function of intensity: there are many more at low
light levels. The STIS 2048 and Tek 2048 #3 have both hot columns and blocked pixels.
These are obvious, and should be avoided. The Thomson CCD has about 20 short traps
visible only at low background levels; at higher backgrounds almost all of these shorten
to a single low pixel.
3.6.4 Standard Stars:
In broadband UBVRI work excellent results have been
achieved by using standards from Landolt (AJ 104, 340, 1992). There are fields (eg T Phe,
SA 98, SA 110) which contain several stars in a 3-4 arcmin diameter field with a wide
range in color, and so are invaluable for use with large CCDs. John Graham's uncrowded
E region standards (PASP, 94, 244, 1982) in the magnitude range of 10 to 16, the older
Landolt standards (AJ, 78, 959, 1973; AJ, 88, 439, 1983) and some of the fainter (10-11
mag) Cousin's stars (J.W. Menzies et al., SAAO Circ., 13, 1) are useful. Copies of these
lists are kept in each of the console rooms. Look carefully at the quoted errors on the
standards before observing them. Residuals in the photometric solutions are a few
percent at worst, and standard stars (properly done) repeat to 0.005 mag or better from
night to night. It is preferable not to mix Landolt standards with those of Cousins/Graham
if you are covering a wide range in color and are interested in getting systematics down
to the 0.01 mag level.
For narrow band imaging one can try using spectrophotometric standards (R. Stone, J.
Baldwin, 1983. MNRAS, 204, 347 and 1984. MNRAS, 206, 241, D. Hayes, 1970. ApJ, 159,
165, M. Hamuy et al., PASP 104, 533, 1992).
The biggest mistake made in observing standards is to underexpose: try for a peak in
the star of at least a few thousand ADU above background (exposure times of 5-20
seconds for V = 12-14 at the 4-m, with the Tek CCDs and gain at 10 e-/adu you can go
a magnitude brighter than these figures). For such short exposures the dominant source
of error is the readout noise in the large measurement aperture necessary to contain
99% of the stellar light. Shutter clocking is accurate enough to permit 1 second
exposures, but be aware of the shutter delay discussed above. If necessary you can
remove the shutter timing effect from the data (Stetson,
1988 IAU reports).
Use of faint photoelectric sequences (particularly those in crowded fields) requires
caution; these often have proven to have systematic errors, and require long exposure
times in each filter to reach the desired accuracy in a large measurement aperture.
The care with which one observes standards will be reflected in the accuracy of one's
photometric calibration. It is common to observe several standards at the start of a
photometric night, several more in the middle, and still more at the end. If you have fields
with only one or two standards, try taking two exposures of a standard field per filter,
stepping the telescope by 20 arcseconds between exposures. The extra frame takes very
little time, and insures against your star landing on a bad pixel or cosmic ray. You could
even use the observe command with focus mode enabled to give you several
images of the standard but only one readout. For the highest accuracy one should
definitely observe enough standards at a range in airmass to solve for the extinction
coefficients. Tololo extinction can be as low as kV = 0.12, or > 0.2 if conditions are dusty,
so one should not rely on the nominal values (in 1991 volcanic dust caused elevated
extinctions). The CCD camera has shown itself capable of doing better than 1%
photometry: it seems a shame to skimp on the external calibration necessary to achieve
this. Furthermore, to make proper use of such an accurate calibration will require the
existence in the object frames of uncrowded stars exposed to the same level as the
standard.
CTIO operates several types of CCDs. There is, as yet, no general purpose CCD
and a CCD suited to your program will be the one scheduled. Loral CCDs
require periodic UV- flooding in order to provide enhanced sensitivity below
5000Å, and must be kept cold in order to retain the improved response. The
output signals from CCDs are extremely small (remember 1 electron = 1.6 x
10^(-19) Coulomb!) and the electronics must not allow significant extra noise
to degrade the performance. In case of problems, electronic components can be
replaced (at the board level) and if necessary, the CCD can be changed. If
you are dissatisfied with the performance of the CCD system, please consult
with Observer Support personnel.
At PFCCD the CCD is mounted in a downward looking liquid nitrogen filled
dewar, and is kept at a constant temperature by a small heater. This
temperature is displayed in the console room and should remain constant. If
the temperature appears to be rising immediately alert either your Observing
Assistant or Observer Support personnel. After initial power on and the dewar
has been filled with liquid nitrogen for the first time, it requires at least
4 hours for all temperatures and operating voltages to stabilize. The dewar
hold time is over 12 hours -filling at the start of the night will generally
do for the whole night, except in winter time when, near midnight, the night
assistant will ask for permission to return the telescope to the access
platform in order to fill the dewar.
At CFCCD the dewar is upward-looking, otherwise procedures are as in the
above paragraph. All the dewars have the fill-tube extending only half-way
into the Liquid Nitrogen can. Although this has the disadvantage that the
can cannot be filled to more than half-capacity, the advantage is that the
dewar can be used in any orientation.
The analog signal (which can be considered to be some number of electrons)
from each pixel of the CCD is first amplified, then passes through an
integrator which discriminates against various noise sources, and is then
converted into a digital signal by an analog-digital- converter (ADC). Hence
the signal at this stage is measured in "counts" or " adu's" (analog-
to-digital-units). The ratio adu/electrons is called the gain, although by
convention we usually talk about the inverse gain, electrons/adu. The gain is
an adjustable parameter on the CTIO CCD systems. The range of data numbers
possible extends from 0 to 65535 (0 to 216-1, 16 bits).
A summary of the
characteristics of the detectors normally used for direct imaging is as
follows:
1. The full well capacity is the limit above which charge "spills"
out of a CCD pixel into adjacent pixels. At some level below full
well the response of the CCD becomes non-linear to incident light.
Some CCDs retain excellent linearity right up to full well, but
for others the departure from linearity becomes severe well below
full well. This will define the practical maximum charge capacity
per pixel. It may depend on details (slope, amplitude, levels) of
the clocking voltages and for some of our devices there is some
hope of improving the values given.
Typically a factor 20 more charge must be accumulated
before noticeable bleeding occurs. This bleeding is stronger in the column
direction for our CCDs. With normal gain settings the data system limit of
65535 counts is reached prior to entering the CCD non-linear region. Care
should be taken always to operate the CCDs in the linear part of their
response.
2. The readout noise is somewhat dependent on gain, larger values of readout
noise correspond to larger values of e-/ADU. Similarly, the read time is
given for gains of 3-4 e-/adu.
4.3.1 The Tek 1024 CCDs
We have one Tek 1024 CCD, #2. It is now (2000) only scheduled in case of
emergency. This is a quad-amplifier device
with very low read noise. Cosmetics are
superb. QE is poor in the UV and U band photometry is not recommended. Note
that these CCDs do not have the serial registers and amplifier areas shielded
from light, consequently on high light level flat field exposures the
overscans show an exponential decay and during reduction should be fitted with
high-order splines. See the software manual for more details.
4.3.2 The Tek (SITe) 2048 CCDs
We have three Tek 2048 CCDs (#3,#5, #6).These are thinned, AR coated devices
like our Tek 1024s. Read noise is very low, and QE is like the Tek 1024's,
except the U band response is much improved.
Tek #3 has several column defects. There are bright columns on each
side and three major column traps near the center. These should all be
avoided. Tek 2048 #6 is the CCD normally scheduled at the 4-m PF, while
#3 is dedicated to the 0.9m, and #5 is used at the Schmidt.
Type "ccdinfo" ro see the characteristics of the CCD you are using.
4.3.3 The Thomson 1024 CCDs
RETIRED!!!!
The observer can control and change several CCD parameters. These are: the
CCD readout format, the binning, the preflash time, and the gain.
Observers should think carefully whether they need all the field (if
not read a ROI) or the resolution (if not, bin 2x2). Even though the
CCDs with quad readout have short read times by big-CCD standards,
substantial gains in efficiency are possible by reducing the format.
4.4.1 CCD Readout Format
A single region-of-interest can be read out, positioned at an arbitrary place
on the CCD. These operate very efficiently, for instance a 1024x1024 ROI
centered on a Tek 2048 takes under 10 seconds to read out, while a 256x256
takes just 2 seconds.
4.4.2 Binning
The ARCON controllers are able to bin pixels in various ways. The readout
noise per pixel remains the same as in the unbinned case. The gain
calibration is unchanged from the unbinned mode. For most direct imaging
applications at CTIO the CCDs tend to undersample the images, so there is
little advantage in binning. The main advantages of binning are that the
readout time is shorter, and less data is created.
4.4.3 Preflash Time
Default is zero. All our CCDs operate without preflash.
(1) Pixel sizes (arc sec)
(2) Field Sizes (arc min per side)
CFCCD
The 0.9m, 1.5m and 4m CFCCD systems have two filter wheels each holding seven
filters up to 3x3 inches square. These wheels are modular, and we have a single
wheel which holds five 4x4 filters. This is useful when an observer brings 4x4
filters for use at these telescopes, or it is desired to use a filter which we
have in 4x4 size and not in 3x3.
SCHMIDT
The Schmidt has a single filter bolt which holds five filters up to 4x4 square.
1. Take a sequence of 2 bias frames. Make sure that the preflash is set to
zero.
2. Take a sequence of 2 dome flats or, better yet, 2 bias frames with a big
preflash value (giving a total signal of approximately 1/2 saturation).
3. Subtract these frames as follows:
BIAS1-2 = BIAS1 - BIAS2 (use imarith)
4. Use the IRAF imstat task to determine MEAN(FLASH1),
MEAN(FLASH2), MEAN(BIAS1), MEAN(BIAS2), RMS(FLASH1-2), and RMS(BIAS1-2). Note
that it is very important that the stats box be positioned on a region of the
CCD free from traps, bad columns, or any bad pixels.
5. Calculate the gain,
e/adu = (MEAN(FLASH1)+MEAN(FLASH2)) -
(MEAN(BIAS1)-MEAN(BIAS2))
RMS(FLASH1-2)2 - RMS(BIAS1-2)2
6. The readout noise in electrons is
RON(e) = RMS(BIAS1-2) * (e/adu) / 1.41
2.2 Telescope Control System
2.3 CFCCD
2.4 HELP!
3. DATA-TAKING PROCEDURES.
3.1. Introduction
3.2 Starting Up, and Some Miscellaneous Advice
3.2.1 Loading & Running the Software.
You should login on the data acquisition computer using your VISITOR username and
password. When you login for the first time, you will be offered a choice of
OPENWINDOWS or SUNTOOLS environments. On subsequent occasions your initial
choice will be loaded automatically. Several windows will open automatically, including
a (blue-colored) IRAF Arcon user-interface window via which you interact with the CCD.
There is also a (brick red-colored) IRAF reduction window, an Arcon console window, an
Arcon status window with countdown timer, and an SAO-image (or XIMTOOL) display
window (in OPENWINDOWS). You will almost certainly want to re-size and move some
the windows around to get a display that you are comfortable with.
3.3 Instrument Setup and Checkout
3.3.1 Introduction.
At the start of your observing run, Observer Support personnel will check that the CCD
is operating correctly, and that all the hardware functions. They will fit your selected filters
into the filter bolts. They will also ask you which gain setting you require. The various
options are listed by typing3.4 Calibration Frames.
3.4.1 Introduction.
Most of the calibration frames can be taken during the afternoon. Almost all observers
prefer to take calibration frames each day, and we also recommend this, but it is not
absolutely essential. Observer Support personnel will help you to get set up if required.
If you have special requirements (dome flats in the morning etc) please consult with
Observer Support in advance so that arrangements can be made. Note that in order to
ventilate the domes so as to minimize "dome seeing" it is normal for the shutters to be
opened approximately one hour before sunset. Please try to finish your calibration frames
prior to this time. This is particularly important for the 0.9-m where the focus is a strong
function of temperature and very poor images can be expected unless the telescope is
at a similar temperature to the outside air3.5 Night Time Procedures.
3.5.1 Start of the night procedures.
About an hour before sunset the dome shutters
will be opened unless calibration frames are still being taken. All doors should be closed
and lights turned off. During twilight, your Observing Assistant will fill the dewar with
liquid nitrogen and check the telescope pointing on a bright star near the zenith. If it is
the first night that the CCD has been installed then the offset between the CCD and the
acquisition TV camera will also be determined, and the focus of the TV camera checked.
Your Observing Assistant is familiar with focussing methods, offset determinations,
operation of the TV and Guider, and will help if required.3.6 Observing.
3.6.1 Finding Your Object:
The quickest procedure is to point the telescope at the
correct position and take a preview or movie. Alternatively the telescope
can be positioned such that the object is on the TV, and then offset back to the CCD.
Then find a guide star.4. DETECTOR CHARACTERISTICS
4.1. General Principles
4.2 Detector Options
Tek 1024 Site 2048
Pixels 1024x1024 2048x2048
Pixel size (microns) 24 24
Readout noise (electrons) 4-6 3-5
Electrons/ADU (typical) 1-4 1-4
Cosmic ray rate (per min) 20 100
U: Sensitivity at the 4m 18 54
B: ( U=B=V=R=I ) 220 288
V: ( = 20 mag ) 350 380
R: ( star, in ) 380 356
I: (electrons/sec) 180 167
Read time (quad, secs) 10 25
Notes:4.3 Detector Notes
4.4 Detector Control Options
4.5 CCD scales at various foci
Tek Tek
1024 2048
1.5-m f/7.5 0.44 0.44
1.5-m f/13.5 0.24 0.24
0.9-m f/13.5 0.40 0.40
Schmidt f/3.5 2.3
Tek Tek
1024 2048
4-m f/7.5 2.70 5.40
1.5-m f/7.5 7.41 14.8
1.5-m f/13.5 4.09 8.18
0.9-m f/13.5 6.76 13.5
Schmidt f/3.5 78
APPENDIX I:FILTERS FOR CCD IMAGING
A
filter list
is available via WWW and will also be found in
NOAO Newsletter
No. 34. A complete CTIO filter list is also available via ftp. See the
WWW home page.
APPENDIX II: GAIN AND READOUT NOISE CALCULATION
The gain in electrons/adu is usually accomplished using the gfind
command (in the contributed package). Here is a cookbook recipe for
determining the gain in electrons/adu, for anyone who wants to do all the
steps one by one.
FLASH1-2 = FLASH1 - FLASH2