Observer's Manual

 

 

Version 1.5 - Aug. '98
by P. Bizenberger
with help from
T. Herbst, R. Lenzen, D. Thompson

Table of Contents

Introduction

    This manual is currently under construction, and likely to change frequently. The authors welcome and suggestions for improvement, additional information which should be included here, and otherwise constructive criticism.

    In this manual, we are assuming that the reader is already familiar with infrared observing in general, as well as the reduction of data into a final, usable form. If not, the reader is directed to the MAGIC manual, where much of this information can be found. We attempt here to provide enough information so the user can prepare for and execute a successful observing program with Omega Cass.
 

S/N Consideration

    In most instances of IR observations, the target object is fainter than the sky background. For a typical exposure with Omega Cass using double-correlated sampling (see below), the read noise Nr is approximately 35 electrons, and the dark current is < 0.1 electrons / second. The photon noise associated with the sky background increases as the square root of time. Therefore, the signal to noise ratio increases approximately linearly with time until sufficient background photons have been collected to dominate the noise sources. This occurs when when the background signal is more than the square of the read noise, or 1225 electrons. With Omega Cass, the conversion factor between electrons and Analog to Digital Counts (ADC - one ADC is the minimum signal unit recorded by the computer) in the image is 4 electrons / ADC. Therefore, any integration with less than approximately 300 sky counts is read noise limited , and the signal to noise ratio increases linearly with time. Longer exposures are background limited, and the signal to noise ratio increases only as the square root of time.

    Integrating for much longer than the time needed to reach the background limit does not provide any advantage to the observer, since a series of independent exposures will also increase the S/N ratio as square root of time. In fact, the variability of the sky level and other factors make it advantageous to end an exposure when the background limit is reached. For broad band images with Omega Cass, this usually requires only a few seconds.
 
 

The Array - Readout Techniques

    Omega Cass's detector is a Rockwell 1024 x 1024 HAWAII array. The detector can be read with the following modes.

Single Correlated Read

    This is the simplest readout scheme. The pixels are reset at the beginning of an exposure, and read out once at the end of the integration. This does not remove the variable pedestal level (kT noise) and any initial offsets which can vary from pixel to pixel. We do not recommend using this mode for observation. Its main usefulness is in checking the signal level for saturation.

Double Correlated Read (Reset Read Read)

    This is the most commonly used mode for general observing. The array is read immediately after the initial reset and again just before the final reset at the end of the integration. This eliminates the kTC noise and other offsets, but increases the read noise by root-2 because the noise from two readouts goes into a single image.

    Minimum integration time is 0.84 seconds but the total time to get an image is 1.68 seconds because of the two required reads.

Full MPIA Mode

    This mode is implemented now. It is supposed to give a higher observing efficiency compared to the existing modes. See the following graph for the calculated efficiency  of a stack of 15 images. The MPIA mode is the common double correlated read mode (Reset Read Read) which is usually used with Omega Cass. The IRL mode is the double correlated read mode used with MAGIC and Omega Prime. For short integration times is the efficiency of the Full MPIA mode significant higher to the other modes. Since the overhead of one mode is constant for a read out stack, the ratio of overhead to integration time is the worst for minimum integration time.

    The noise is the same as for the Double Correlated Read. Minimum integration time is 1.68 seconds.

    The common double correlated read out modes are 'frame orientated' i.e. you reset the hole array, read the hole array, read again the hole array and subtract the two frames. Assuming a very fast reset, this takes twice the time to read the array, to achieve an image with minimum integration time. In this case, a single pixel integrates light only as long as it takes to read the array once. The resulting efficiency is for minimum integration time 50%, changing to better values for longer integration times.

    The new Full MPIA mode is 'line orientated'. You read one line, reset the same line and read it again. Do this for the hole array. Next cycle is the same. You read one line, reset the same line and read it again. To archive a double correlated image, you subtract the second read of the first cycle from the first read of the second cycle. Repeat this for the hole array. The efficiency is in this case almost 100%, it is not exact 100% since the reset is not infinite short and the first read of the first cycle (as well as the second read of the last cycle) is lost and counts as an overhead. This overhead becomes negligible when taking many frames in a row. The graph is for a stack of 15 images (Repeat 15). For a single image, this mode shows no advantage to the MPIA mode, it rather has the disadvantage of the longer minimum itegration time.

        The minimum integration time increases from 0.84 to 1.68 seconds because still two reads are necessary for a double correlated read. But the integration of photons is done during the hole read process.

Subarrays

    This readout scheme reads a square box of variable size.  This mode is well-suited to occultations. The actual readout technique is Reset Read Read. The size of the subframe must be an even number. In addition the subarray has to be located entirely in one quadrant of the detector.

    The readout speed depends on the size of the subframe (see table below).

size of subarray
/ pixel
# of frames
total time
for all frames
/ sec
min. integration time
per   frame
/ msec
16 x 16
1000 
  2.2
 1.0
"
10000 
21.2 
1.0
32 x 32
1000
  7.3
 4.0
"
3000
21.3 
4.0
64 x 64
1000
 27.4 
 14.0 
data for the position described in the following

    The minimum integration time for a subarray depends of the location on the array. The pixels of a quadrant are clocked line by line. To read a subarray, all previous lines to the subarray must be clocked as well as all previous pixel in a line where the subarray is located. To avoid an overhead by clocking unused lines and pixels, the best location for a subarry is in the corner of an quadrant where the clocking of a quadrant starts. In the following image, the quadrants are labeled #1 to #4 and the direction of clocking is indicated. Taking in account that the optics have less aberrations in the center region, it is best to locate the subarray in the lower left corner of quadrant #4. At this position, the best performance for subarrays is archived. See the square, orange box.

    The best position for subarrays using ALFA is the lower right corner of quadrant #1. This is due to an additional mirror of the image by the ALFA optics.

The Camera

Optical Set-up

    Omega Cass is a re-imaging camera with three different image scales accepting input beams up to f/8 . The scales get changed by changing the cameras of the optic which are mounted on a wheel remotely controlled by software. The total length of the optical path is constant for all three scales.

    All filters, grisms and polarisiers are close to the pupil position in the parallel beam. The optics are designed for the whole spectral range from 1.0 to 2.5 µm. There is no refocusing necessary due to a filter change or optic change.

Mechanical Set-up

    The Omega Cass dewar is equipped with two tanks for liquid nitrogen. The outer tank and the connected outer baffle operates as a radiation shield. The inner tank and inner baffle preserves an area stabilized at 77 Kelvin. Here, all the optics, filters and masks are mounted. The detector is direct mounted to the cold plate of the inner tank. There are four temperature sensors mounted in the dewar. The temperature should be:
name of sensor
 nominal temperature
location
outer
~ 90 K
mounted on top of the outer shield
inner
<80K
mounted on top of the inner shield
filter
<80 K
mounted on the filter box
array
<80 K
mounted next to the detector

  For all temperature sensors is only one monitor available, to display all sensors switching cables is necessary.

Data Handling

     It takes a considerable amount of time to transfer the data from the camera and save it to the hard-drive on the workstation. To reclaim some of this otherwise lost time, the software has been configured with two image buffers. Thus, a new image can be read out while the previous image is being saved. (See also the Software manual, section 11 Macro Format)
 

Mounting and Aligning

    The camera must be aligned with the telescope before each observing run!

    Usually, the camera will be mounted and aligned by the Calar Alto staff but the observer must double check to guarantee a perfect set-up. Aligning will take place in the beginning of the first night of observing since you need a star to do this. To align the camera with the telescope optical axis is important in terms of background reduction. The entrance pupil of the telescope must be aligned with the Lyot stop of the camera. If you are not familiar with this procedure please see the MAGIC manual for basic instructions.

    To align the entrance pupil of the telescope with the Lyot stop of Omega Cass, you have to apply two different methods for the two directions (north-south and east-west). The mounting flange can be tilted in one direction which allows the alignment in the east-west direction. The north-south direction can only be aligned by rotation of the Lyot stop wheel.

    Aligning a single wheel can be done with the following command    wheel # rel xxxx    in the interpreter window of the software.
Where # is the number of the wheel (0 = optic wheel, 1 = grism wheel, 2 = lyot wheel, 3 = filter2 wheel, 4 = filter1 wheel, 5 = mask wheel) and xxxx is the number of steps. All wheels have a different gear ratio i.e. a given number of steps moves the wheels a different amount.

    If you have aligned a wheel, you must edit the corresponding info file with the new position. Otherwise the software will remember the initial position and will use it again for further movements. After updating the info file, you also must tell the software that the info files have changed. Apply the command wheel rdb and this will update the software automatically.

Note! The info files are located in a  privileged area. A password is required to allow access, ask the night assistant for help.
 

Observing with Omega Cass

Startup Procedure

    The first day of your observing run you will get an introduction to the telescope, camera and software from the Calar Alto staff. To start up the software see the Software Manual.

Direct Imaging

    The optical design of Omega Cass offers three different image scales by exchanging individual cameras while the collimator is the same for all systems.

    Omega Cass offers you the following pixelscales for the following telescope configurations:
 

3.5 m - f/10  ~ 0.3 "/pixel 
~ 0.2 "/pixel 
~ 0.1 "/pixel 
3.5 m - with ALFA 
              f/25 		 ~ 0.12 "/pixel 
~ 0.08 "/pixel 
~ 0.04 "/pixel 
3.5 m - f/45 ~ 0.067 "/pixel
~ 0.044 "/pixel 
~ 0.022 "/pixel 
2.2 m - f/8  ~ 0.6 "/pixel 
~ 0.4 "/pixel 
~ 0.2 "/pixel 

    Omega Cass offers a set of broad band and narrow band filters for direct imaging. See the Technical Characteristics for a full list of filters.

Spectroscopy

    NIR-spectroscopy can be done using Omega Cass in long-slit or slitless mode:

    At present, two grisms are available for Omega Cass:

    The calibration curves for grism3 and grism4 are available.

    Two additional grisms are in preparation:

    Long slits of 24 mm length are offered, corresponding to 2.36 arcmin (f/10) or 0.95 arcmin (f/25) with different slit widths. See the  Technical Characteristics for the current set-up of the mask wheel.

    Objects can be centered into the slit by first taking a direct image for the selected wavelength band. The image of the object can now be moved to the required position on the detector. Now the slit is switched in, a direct image
is taken through the slit. Recentering is provided if necessary. For a last step, the required grism is positioned.

    Wavelength calibration can be done by using sky spectrum or by switching in a calibration lamp (Argon). The argon spectrum is given in the following wavelength ranges: 1.0 - 1.25 µm, 1.25 - 1.50 µm, 1.50 - 2.0 µm,  2.0 - 2.5 µm and 2.5 - 3.5 µm. The last range is not used for Omega Cass of course but it is for completeness.

    In principle, the spectroscopic mode of Omega-Cass can be combined with the adaptive optics system ALFA, however, using this combination no slit rotation at the sky will be possible, the fixed position angle is -14 deg. A rotator, that will provide slit rotation even in combination with ALFA is in preparation.
 

Polarimetry

    There are two polarimetric modes for Omega-Cass: For very compact or point like objects the double-imaging or Wollaston-mode is used: Two Wollaston prisms are providing double images of about 17 arcsec beam separation (f/10), one in North-South direction, a second one for a position angle of 45 deg. Wollaston prisms are perfect polarizers in the 1-2.5 µm region. The beam separation depends of the wavelength. A calibration curve is available

    The second mode uses wire-grid polarizers which provide single linear polarized images. Four single images have to be taken through the four offered analyzers which are mounted at position angles of 0, 45, 90 and 135 deg.
 

Calibration

Sky Subtraction

Point Sources-Moving Sky

When the observing targets are pointlike or small, most of the array is, in fact, measuring the sky background level. Offsetting the telescope between exposures produces a sequence in which a subsection of the array contains the target object in one exposure and sky measurements in all the others. A median average of these sky frames will eliminate the pointlike sources (i.e. an image in which each pixel is replaced by the median of values at that location in all the sky frames) . Taking the median of, say, four neighbouring frames will give a sky level that is well-correlated temporally and spatially with that in the target exposure. This so-called "moving sky" frame works extremely well with virtually no penalty in observing time, since you should already be moving the telescope slightly between frames to account for bad pixels. You must move the telescope by at least twice the size of the largest observing target for this technique to be effective.

Extended Sources

When the observing target is extended over a significant portion of the detector, the moving median will be contaminated by source flux and will not reflect the current sky level. In these circumstances, the observer must offset the telescope to a relatively source-free area to measure the sky. A typical sequence may be Object-Sky-Sky-Object-Object-Sky-Sky... The atmospheric conditions are the major factor in deciding how frequently to measure the sky level. Note also that most "empty" fields contain numerous stars, particularly given the large field and high sensitivity of Omega Cass. You should offset the telescope between sky exposures and form the median as before to eliminate point sources. The positive aspect of taking frequent off-source exposures is that it gives you an opportunity to monitor the atmospheric transparency, seeing, and focus using a star in the sky field.
 

Flat Fielding

The flat field image provides a means of removing pixel to pixel gain variations in the array. There are several different techniques for producing the flat field, but they all share the goal of exposing the detector to uniform illumination. The image values in the resulting frame will be proportional to the pixel gains. Division of the object frames by the flat field will eliminate these variations.
 
 

Linearity

The calibration procedures assume that the voltage on the detector is linearly proportional to the incident flux.

A plot of the relationship fot the detector in Omega Cass will be placed here. (soon)

Note that the linearity plot uses the median counts for all pixels. Some pixels are more nonlinear and others less. Observers who need accurate photometry will want to correct all their exposures for nonlinearity before proceeding with the standard data reduction. We recommend using a second or third order polynomial fit to each detector's response.

A linearization matrix is not yet available.
 

Flux Zero Points

Calculated Counts for Zero Magintude according to the formula:      CFZM = (source - sky) *  10^ (0.4 mag)
 
 
Telescope and Camera Configuration J Filter H Filter K' Filter K Filter
3.5 m f/10 - 0.3 arcsec/pixel
        
3.5 m f/10 - 0.2 arcsec/pixel        
3.5 m f/10 - 0.1 arcsec/pixel
        
3.5 m f/25 - 0.3 arcsec/pixel
        
3.5 m f/25 - 0.2 arcsec/pixel        
3.5 m f/25 - 0.1 arcsec/pixel
        
 !! Waiting for observer's feedback !!

Faint IR Standard Stars

    There is a list of UKIRT faint standard stars from Casali (1992 UKIRT Newsletter, 4, 33) available. Dave Thompson prepared the corresponding finding charts and also a list of interpolated K' magnitudes.
 

Troubleshooting

     If you get images like this, check the entrance window of the dewar if water is condensed.
 


Peter Bizenberger
biz@mpia-hd.mpg.de
Tel. (+49) 6221 528311