PDS_VERSION_ID = PDS3 OBJECT = INSTRUMENT INSTRUMENT_HOST_ID = MSSSO INSTRUMENT_ID = CASPIR OBJECT = INSTRUMENT_INFORMATION INSTRUMENT_NAME = "CRYOGENIC ARRAY SPECTROMETER/IMAGER" INSTRUMENT_TYPE = "INFRARED IMAGING SPECTROMETER" INSTRUMENT_DESC = " Instrument Overview =================== This manual describes the operation of CASPIR, the Cryogenic Array Spectrometer/Imager on the ANU 2.3 m telescope at Siding Spring Observatory. CASPIR uses a Santa Barbara Research Center (SBRC) CRC463 256 x 256 InSb detector array to provide direct imaging and spectroscopic capabilities in the 1-5 micrometer wavelength range. Two direct imaging focal plane scales of 0.5 arcsec/pixel and 0.25 arcsec/pixel are available, as well as long slit J, H, and K grisms giving two pixel resolving powers of ~500 through a 1 arcsec x 128 arcsec slit, and IJ, JH, and HK cross-dispersed grisms giving two pixel resolving powers of ~1100 through a 1 arcsec x 15 arcsec slit. Coronograph and imaging polarimetry functions are also available. Operational Considerations ========================== CASPIR operates within the environment common to all infrared instrumentation on the 2.3 m telescope. All mechanical functions are controlled from MOPRA to the LSI-11/23 located in the Cassegrain Instrumentation Rack in the Nasmyth Lab, and from this to the instrument control subrack mounted on the Instrument Mounting Box (IMB). The detector array clocks, biases, and signal processing are performed by the SBRC Array Control Electronics (ACE2) also mounted on the IMB, close to the CASPIR dewar. The ACE2 is controlled directly by MOPRA through a 9600 baud RS-232 connection. Data from the four detector output channels are digitised in the ACE2 using four 16-bit, 500 kHz Burr-Brown ADCS, and serialised using four transputer Link Adaptors. The serial data are transmitted from the Cassegrain focus to the Nasmyth Lab where four T800 transputers receive the data and process it as necessary. When the requested integration sequence completes, the data are transferred to MOPRA through a transputer link to Q-bus interface and are displayed on the workstation screen and stored on disk. The detector array temperature is controlled by a commercial controller located in the Nasmyth Lab. MOPRA communicates with this controller through an RS-232 line. Optics ====== CASPIR is a cryogenic reimaging camera with a 50 mm long, 10.4 mm diameter collimated beam section. The camera body is cooled to ~60 K by the first stage of a closed cycle helium refrigerator, and the detector array cooled detector array cooled to ~32 K by the second stage of the cooler. The dewar incorporates a novel design which uses five 16-position annular wheels mounted coaxially around the cooler to produce a compact vacuum system. The wheels are driven by motors located on the dewar base plate. The CASPIR dewar mounts on port A of the IMB. The rotatable dichroic mirror in the IMB directs the f/18 telescope beam to the dewar. The dewar window is a Sapphire/CaF2 doublet which acts as field lens to image the telescope exit pupil (the secondary mirror) onto an internal cold stop. A cold gold-coated mirror then directs the beam down, parallel to the dewar axis. The telescope focus is located immediately below this mirror at the Aperture Wheel. The Aperture Wheel contains baffles for the 0.5 arcsec/pixel and 0.25 arcsec/pixel focal plane scales, a range of slits for the grisms, coronograph masks, and the field mask used for imaging polarimetry (Table 12). The diverging beam then passes to a fixed MgO/CaF2 doublet, collimator lens which produces the collimated beam section. Immediately below this, in the collimated beam, is the Upper Filter Wheel which contains the filters listed in Table 13. Next is the Utility Wheel which is located at the pupil plane. This wheel contains the direct imaging cold stop, the MgO/CaF2 doublet slow camera lens for the 0.25 arcsec/pixel focal plane scale, the six grisms, and the Wollaston prism polarimeter analysers (Table 14). Below the Utility Wheel is the Lower Filter Wheel which contains the filters listed in Table 15. Note that some of the broadband filters require blocking filters located in the Lower Filter Wheel. Both filter wheels contain clear positions. The detector array should not be exposed to optical light while cold, so the software prevents both clear positions being selected at the same time. Note also that the Lower Filter Wheel is located in an f/10.4 converging beam when the slow camera is used, so some refocusing between filters may be required. The final wheel is the Lens Wheel which contains the MgO/BaF2 fast camera lenses for the 0.5 arcsec/pixel focal plane scale. These are rotated out of the beam when the slow camera is used. The detector array is located at the lower end of the dewar at the camera focus. Detectors ========= The detector is an SBRC CRC463 256x256 InSb array which is sensitive from ~0.9 micrometers to 5.5 micrometers. The array has four output channels corresponding to four interlaced columns (12341234...) The array is controlled by the SBRC ACE2 drive electronics which is mounted close to the CASPIR dewar. Communications with the ACE2 is through an RS-232 connection to MOPRA. The CRC463 is a hybrid device in which the InSb detector material is bump-bonded to a silicon multiplexer through indium bumps. The multiplexer is a switched FET read-out device, which operates differently than a CCD. Each pixel (or unit cell) contains four FETS; the two row select FETs and the reset FET are switches which can be thought of as closed when activated. The fourth FET acts as a source-follower amplifier which continuously samples the voltage on the detector node without affecting its value. A load is provided for the unit cell source-follower FET. This load FET is located in the column biasing circuitry. The output FET acts as a second source-follower amplifier with its external 10 KOhm load resistor in the ACE2 electronics rack. The two column select FETs also act as switches. Pixels in the array are sequentially addressed by pulsing column and row shift registers which activate the column and row select FET switches. Once a pixel is selected, the voltage of the detector node can be non-destructively read via the source-follower amplifier signal train, and the detector node voltage can then be optionally reset to Vdduc by activating the reset FET switch by pulsing the Phibrst clock line. Note that the stored charge is not transferred across the array like a CCD. Instead each pixel is sequentially reset and read. This results in a time delay in the integration window across the array of one frame readout time between the first and last pixels. The array readout scheme permits the use of a variety of readout methods which are now described. Data Modes ========== Readout Method 1: Fast Sampling In the fastest readout method, speed is considered more important than accuracy so we take only one sample per pixel and reference this voltage to electrical ground (Vss). The detector node voltage is applicable in readout method I. The unit cell is reset, integrated, and sampled once at the end of the integration ramp. This method should be used for imaging in the 3-4 micrometer band and at M where the high thermal background flux significantly fills the detector wells in of order the frame readout time. In this situation, the dominant noise source is photon shot noise from the background flux, so readout noise is not an issue. The minimum readout time for this method is 0.2 sec. Readout Method 2: Relative Sampling Under less extreme background conditions, significant improvement in stability can be made by referencing the signal level to the reset level, instead of electrical ground. This is done in readout method 2. Note that the detector node voltage jumps when the reset FET is switched off, and this pedestal level is not removed in readout method 2. The pedestal level is different for each pixel in the array so this imprints a pedestal structure on the image, which is difficult to remove completely. The uncertainty in the value of this pedestal is known as kTC noise. For direct imaging through broad band filters with high sky levels, the dominant noise source is still the photon shot noise of the background flux so readout method 2 gives acceptable performance. The minimum readout time for this method is 0.3 sec. Readout Method 3: Double-Correlated Sampling The pedestal structure can be removed using a readout method which references the detector node voltage at the end of the integration to the voltage at the beginning of the integration. In principle, this readout method is susceptible to electrical 1/f noise since it differences two samples separated in time by the duration of the integration. In practice, this is unlikely to be an important noise source. This is the preferred readout method for broad band imaging because it does not imprint the pedestal pattern on the data. The minimum readout time for this method is 0.4 sec. Readout Method 4: Triple Correlated Sampling The potential 1/f noise problem with readout method 3 can be overcome, at the expense of two more reads, by referring both the start and end reads to their respective reset levels. Naively, this should be the most accurate readout method to adopt. However, we must now delve into the more obscure operating characteristics of the CRC463 array to see why other effects dominate. Readout Method 5: Fowler Sampling The FET switches in the CRC463 multiplexer are not perfect switches as assumed above, but instead have finite gate capacitances that act as sinks of charge that would otherwise remain on the detector node capacitance. The reset pedestal (i.e., the amount the detector node voltage jumps by when the reset is taken off) is due to a redistribution of charge from the detector node capacitance to the reset FET gate capacitance that occurs when the reset FET gate voltage (i.e., the reset clock voltage) moves positive to switch the reset FET off. This amounts to ~100 mV of lost detector reverse bias (or well depth). Similarly, a further ~400 mV of detector bias is lost when the row and column select FETs are switched off to deselect the pixel. This constitutes a movement of ~500 mV of charge off the detector node compared to the normal operating detector reverse bias that remains of only ~200 mV. The true pedestal is therefore significantly larger than the reset pedestal seen if each pixel is reset and read on one pass through the array. Fowler & Gatley (1990, ApJ, 353, L33) show that the read noise can be reduced by performing multiple nondestructive passes through the array at the beginning and end of the integration ramp. By resetting each pixel on one pass through the array, and sampling the detector node voltage on subsequent passes through the array, the true pedestal is removed from the data. Each time a pixel is selected charge is redistributed from the row and column select FETs back onto the detector node capacitance. Fowler claims that the read noise is predominantly due to the kTC noise associated with this charge redistibution. By performing multiple non-destructive passes through the array at the beginning and end of the integration, the read noise is reduced by the square root of the number of passes. Readout Method 5 implements Fowler sampling in this way. The number of reads at each end of the integration is set by the FNDR parameter (CASPIR/FNDR=...). For applications where low read noise is required, at the expense of increased frame readout time, method 5 is the preferred readout method. This is likely to be the case when using the grisms. Processing ========== Well Depth Considerations The detector node capacitance for the CRC463 array is ~0.06 pF. The well depth is then proportional to the actual reverse bias voltage across the detector; q = CV. Consequently, a detector reverse bias of ~160 mV is required for a well depth of 60,000 e. Our array has an odd-even column effect which causes the applied detector reverse bias on even columns to be lower than that on odd columns. This means that even columns have well depths ~80% smaller than odd columns. Linearity Correction When the bias-subtracted linearity data are plotted as signal rates, it is apparent that the CASPIR array has a quadratic non-linearity that must be allowed for during data reduction. The best description of this non-linearity is currently given by the equation: Linear Counts = Raw Counts + 6.4 * 10^-6 * (Raw Counts)^2 It is possible that different quadratic coefficients apply to different pixels. Further investigation of this effect is required. Transputer Preprocessor CASPIR is designed to operate in the high background conditions encountered at long wavelengths. It uses four 16-bit 500 kHz analog-to-digital converters to digitize the data, with the rest of the data train being capable of sustained data rates of at least 2 microsec/pixel. The requirement that individual frames be coadded at this data rate to build up the image means that the data cannot be input directly into MOPRA. Instead a transputer-based preprocessor is used. A transputer is a fast microprocessor chip with considerable in-built parallelism and which uses four fast serial 'links' for I/0. Each link can be connected to another transputer, or to an external device through a 'Link Adaptor', which is essentially a bidirectional serial to parallel converter. This architecture has made transputers popular in parallel computing applications. The CASPIR system uses four transputer link adaptors on the ADC cards to serialise the data at the Cassegrain focus. Four serial lines then bring the signals to the Nasmyth Lab where the transputer preprocessor is located. This consists of four T800 transputer boards which each have 1 Mbyte of memory and each are responsible for processing the data from one serial line. The current frame is DMAD into transputer memory at the same time that the previous frame is being coadded to an accumulation array where the summed image is stored. After the requested number of coadd cycles, the accumulated image is divided by the number of coadd cycles, optionally has a bias frame subtracted and is divided by flatfield frame, and is copied to MOPRA via a transputer link to Q-bus interface with the four separate data channels correctly interlaced to form the final image. The result of each sequence of coadd cycles is displayed on MOPRA's workstation screen and is stored on MOPRA's disk. The individual frames from each cycle are not normally saved (this only occurs in the occultation observing mode). Mounting ======== Instrument Mounting Box All infrared dewars mount on a box attached to the Cassegrain focus of the 2.3 m telescope that is known as the Instrument Mounting Box (IMB). Two dewars can mount on the IMB at the same time and a dichroic mirror in the IMB can be positioned to direct infrared light from the telescope to either dewar. The IMB output ports are identified by the labels A-D. From the Cassegrain Access Platform, with the control electronics racks on your left, the port position facing you is 'A', the one to your right (not used) is 'B', the back position is 'C', and the control electronics mount on face 'D'. A dichroic reflexer directs the infrared beam to the selected IMB port while an acquisition system views the same field in optical light. Auxiliary modules can be mounted above the dichroic reflexer in the IMB. A calibration lamp module is available, and a polarimetry module is under construction. A manual dust cover is located in the mounting flange at the top of the IMB. The open and close positions are clearly marked on the flange, and it should be used. Calibrations ============ Calibration Module The calibration lamp module can be used to obtain wavelength calibration arc spectra. The module contains Xenon and Argon lamps as well as an incandescent lamp which may be useful for flatfielding. A rotary mirror is used to select one of the three lamps and a flip mirror is placed in the telescope beam to direct the lamp light to the detector. Polarimetry Module The polarimetry module contains an achromatic half-wave plate borrowed from the polarimetry module used at the AAT with IRIS, and a 50 mm diameter wire grid for position angle and efficiency calibration. The warm half-wave plate is rotated in the infrared beam to rotate the intrinsic plane of polarization of an astronomical source with respect to a cold analyser mounted within CASPIR. The analyser is either a Wollaston prism mounted in the Utility wheel, or a wire grid mounted in the Upper Filter wheel. The Wollaston prism separates the parallel and perpendicular polarization components equally on the array, so is a dual beam system. The wire passes only one polarization, but it does allow the grisms mounted in the Utility wheel to be used for single beam spectropolarimetry. Field of View ============= Sky Subtraction The strong and variable near-infrared background has contributions from OH airglow in the J, H, and K bands, moonlight (either directly or reflected off clouds) especially in the J band, and from thermal emission from the telescope and sky in the K and L bands which varies with temperature and humidity. Although the 10-30% variations in background caused by these factors do not strongly limit the S/N of observations (except at K and L for large changes in temperature), they greatly complicate both the creation of mosaics of large regions and accurate surface photometry of objects with extents comparable to CASPIR's field of view. For such observing programs, it is best to obtain sufficient object exposures (and intermixed sky exposures if necessary) to create a SKY frame for each dataset. For programs with single or a few observations of many objects, a sky calibration based on observations of several objects, possibly combined with subtracting a fitted surface from the final image, is the best that can be accomplished. These grouped observations could be treated as one dataset for the purposes of sky subtraction. It is useful to remember that variable airglow can cause the sky background to vary at H by a factor of 2 and at J by 40% on hour timescales. SKY frames for a dataset are created using the redimage task by setting the subsky flag, and supplying values to the obstype, subtype, nrun, and destripe parameters, obstype defines the type of sky observations in the dataset. obstype=all indicates that all images in the dataset are to be included in the creation of SKY frames. obstype=objskyobj indicates that the first image in the dataset is an object image, and this is followed by a sequence of an off-source sky image and an object image, ending with an object image. Only the off-source sky images will be included in the creation of SKY frames. obstype=skyobjsky indicates that the first image in the dataset is an off-source sky image, and this is followed by object and off-source sky image pairs, ending with a sky image. Only the off-source sky images will be included in the creation of SKY frames. obstype=skyobjobjsky indicates that the dataset consists of sequences of sky, object, object, sky frames. Only the off-source sky images will be included in the creation of SKY frames. obstype=radio is a special pattern used for nbL band observations, initially of radio galaxies. Sets of ten dithered nbL frames are separated by Kn sky and object frames that (hopefully) allow drifts in telescope pointing to be corrected in the final mosaicing. It is likely that, these patterns will include most observing sequences in user defined DO files. redimage can be extended to include other sky types if this proves necessary. Standard star measurements recorded in pairs with the star displaced on the array can be processed by selecting obstype=standard. This is a special type requiring exactly two input images. An output image is formed by subtracting the second (sky) image of the standard star from the first (object) image of the standard star. A permanent output file is produced with the name stdnnn_mmm where nnn is the number of the first standard star image and mmm is the number of the second standard star image. This sky- subtracted standard star image can be automatically processed in each of the following steps except that mosaicing and creating a coordinate grid will be ignored. These steps are not applicable to this image. It is most likely that users will fix bad pixels and then measure aperture photometry on the standard star image after sky subtracting. The subtype parameter in redimage defines the type of sky subtraction that is performed. subtype=all defines that all sky images in the dataset will be included in the creation of a single SKY frame, which is then scaled to the median pixel value of each object image and subtracted from them. This is adequate for small datasets where the total time span of the observation is less than about 20 minutes. Larger datasets need to be subdivided into smaller units, with individual SKY frames. This is achieved by setting subtype=running. This causes a SKY frame to be formed for each object image in the dataset from the median of nrun sky images taken immediately before and after the object image. The object image itself is not included in the running median. The SKY frame created for each object image is then scaled to the median pixel value of the object image and subtracted from it. The destripe parameter in redimage determines whether a residual column bias pattern is to be defined and subtracted from each image after normal sky subtraction. Usually this will not be necessary. However, nbL images obtained with readout method 1 suffer from DC drifts in the bias levels of the four output amplifiers between the object and sky frames that are manifest as a residual column bias pattern with four pixel period that is often not removed by normal sky subtraction. When the destripe parameter is set, redimage determines the shape of this bias pattern by projecting the image in the column direction to a 1D spectrum, and then subtracting this spectrum off each row in the image. System Performance ================== The best empirical estimate of the system performance is the observation that for imaging with 0.5 arcsec pixels after 6 hr of on-source integration, objects spread over ~2 arcsec with Kn ~19.5 mag are detectable with a signal- to-noise ratio of ~5, depending on how they are measured. This is confirmed by similar observations with 4 min on-source integration times reaching Kn ~17.0 mag with similar signal-to-noise ratio. The following describes measurements of basic system parameters, and then calculations of the theoretical performance based on these parameters. These calculations can be used to estimate system performance in other configurations, but should be normalized by comparison with the above observed sensitivities. System zero point offsets are based on the total ADUs in a sky-subtracted stellar image after correction for airmass effects and are defined by the equation ZP = M_std + 2.5*log(ADU/sec). Typical values of the zero point offsets for each filter (mainly with the 0.5 arcsec pixel scale) are listed in Table 3. These can be used to calculate the total signal expected on an object of a given brightness or the signal/pixel on an object of a given surface brightness. Table 3: Typical Zero Points Filter Zero Point -------------------------------------- J 21.6 H 21.6 K 20.8 K' 20.6 Kn 20.5 M 11.7 MSO [Fe II] 18.6 H_2 1-0 S(1) 18.0 H I Br_gamma 17.6 2.21 micrometer Continuum 19.1 CO (delta_v=2) 18.5 3.28 micrometer Dust 16.6 3.60 micrometer Continuum 17.3 4.00 micrometer Continuum 16.6 H I Br_alpha 16.3 Typical background brightnesses measured with CASPIR are listed in Table 4. The expected background photon fluxes can be calculated from the tabulated background brightnesses and the system zero point offsets. Table 4: Background Brightnesses (mag/arcsec^2) Filter IRPS IRIS CASPIR CASPIR PICNIC 1984 1993 Nov 1993 Mar 1993 Aug 1994 --------------------------------------------------------------------------- J 15.5 15.0 18.7 .... .... H 14.5 13.7 14.6 .... .... K 11.5 12.5 11.6 11.7 .... K' .... 13.7 12.5 12.6 .... Kn .... 13.2 12.4 12.5 13.1-13.9 2.21 micrometer Continuum .... .... 12.9 .... .... CO (delta_v=2) .... .... 10.7 .... .... 3.28 micrometer Dust .... .... 3.1 .... .... 3.60 micrometer Continuum .... .... 3.8 .... .... 4.00 micrometer Continuum .... .... 2.6 .... .... H I Br_alpha .... .... 2.0 .... .... The noise in an image is a combination of photon shot noise from the sky and telescope, photon shot noise from the object, shot noise from the dark current, read noise, and other systematic noise sources that are difficult to quantify. For small signals, the noise per pixel can be estimated from the equation: Noise = [RN^2 + T*(i_b + i_d)]^1/2 where RN is the read noise in e-, T is the integration time in sec, i_b is the background signal in e-/sec, and i_d is the dark current in e-/sec. The read noise for the double sample readout methods (methods 2-4, and method 5 with FNDR=1) is ~50 e-. The dark current is typically 10 e-/s/pixel for most of the array, but there are a significant number of detectors with dark currents of 100 e-/s/pixel, and some with 1000 e-/s/pixel. With these data, theoretical performance figures for CASPIR can be calculated from the measured background brightness and the camera throughput as quantified by the system zero point. For example, for a Kn background brightness of 12.4 mag/arcsec^2 and a Kn zero point offset of 20.5 mag, the Kn background flux for 0.5 arcsec pixels is (10^[(Z*P-m_std)/2.5])/4 = 434 ADU/sec/pixel In a 5 sec integration, the background count is ~2170 ADU/pixel, or 19550 e-/pixel (1 ADU = 9 e-). The shot noise of this background signal is (19550)^(1/2) = 140 e- which dominates the typical readout noise of 40-60 e-. Consequently, Kn images with 0.5 arcsec pixels and an integration time of 5 sec should be background limited. The total noise per cycle is ~(50^2 + 19550)^.5 ~148 e-/pixel or 148/9 = 16.5. ADU/pixel which is reduced to 4.8 ADU/pixel when 12 cycles are averaged. The measured value is ~5 ADU/pixel. We assume here that sufficient sky frames are averaged so that sky subtraction is essentially noiseless. For a 5*sigma detection of an object spread over n x n pixels, we require an average signal in each of these n^2 pixels of five times the noise per pixel. The total object signal is then n^2 x 5 x 4.8 ADU. This can be converted to a Kn magnitude after dividing by the integration time of 5 sec and using the Kn zero point offset of 20.5 mag. In this way, we can estimate limiting magnitudes at Kn for a range of seeing or object sizes. The results of these calculations are shown in detail in Table 5. Table 5: Performance at Kn (0.5 arcsec/pixel) Image Size 5*sigma Total Signal Mag Time to 20 mag (arcsec) (pixels) (ADU/5 sec) (1 min, 5:1) (min) ------------------------------------------------------------------------ 1 x 1 2 x 2 96 17.3 145 1.5 x 1.5 3 x 3 215 16.4 760 2 x 2 4 x 4 382 15.8 2290 5 x 5 10 x 10 2384 13.8 91200 Predicted performance figures for 5*sigma detections in 1 min of on-source integration in different seeing conditions for various filters are listed in Table 6 for the 0.5 arcsec/pixel scale and in Table 7 for the 0.25 arcsec/pixel scale. The 1 min integration time does not include time for sky measurements and the dead time between frames of ~20 sec. It is recommended to limit individual frames to 1 min exposures, effectively making the elapsed time ~80 sec, so that an adequate number of sky frames can be obtained in the time scale of 15 min on which the sky level is observed to change significantly. If off-source sky measurements are necessary, it is recommended that equal time be spent on the object, and sky positions. Relative performance figures for each filter are of interest in deciding which passband is most sensitive for a particular observation. These calculations for the 0.5 arcsec/pixel scale in 1.5 arcsec seeing and an integration time of 5 sec with 12 cycles are listed in Table 8 for a 15.0 mag star with zero color (S/N_HotStar), a typical unreddened late-type star with K = 15.0 mag, J - K = 1.0, and H - K = 0.2 (S/N_CoolStar), and a typical AGN power law spectrum (S/N_AGN). Performance figures for the grisms can be estimated assuming a slit transmission Tau_Slit = 0.5 and a grism transmission Tau_Grism = 0.5. For example, consider an observation using the HK grism and a 1 arcsec (2 pixel) wide slit of an object spread over 2 arcsec (n_y = 4 pixels) along the slit and recorded in two frames. Table 6: System Performance (0.5 arcsec/pixel) Filter Time Cycles RN^2 Background Dark 5*sigma, 1 min Magnitude Signal Signal 1 x 1 1.5 x 1.5 2 x 2 5 x 5 (sec) (e-)^2 (e-) (e-) (arcsec) ----------------------------------------------------------------------------- J 5.0 12 2500 3100 50 19.1 18.3 17.6 15.6 H 5.0 12 2500 7100 50 18.8 18.0 17.3 15.4 K 5.0 12 2500 53850 50 17.1 16.2 15.6 13.6 K' 5.0 12 2500 19550 50 17.4 16.5 15.9 13.9 Kn 5.0 12 2500 19550 50 17.3 16.4 15.8 13.8 [Fe II] 10.0 6 2500 900 100 16.8 15.9 15.3 13.3 H2O 10.0 6 2500 3810 100 15.8 15.0 14.3 12.3 Cont2.2 10.0 6 2500 6800 100 16.7 15.9 15.2 13.2 CO 10.0 6 2500 29660 100 15.5 14.6 14.0 12.0 Table 7: System Performance (0.25 arcsec/pixel) Filter Time Cycles RN^2 Background Dark 5*sigma, 1 min Magnitude Signal Signal 1 x 1 1.5 x 1.5 2 x 2 5 x 5 (sec) (e-)^2 (e-) (e-) (arcsec) ---------------------------------------------------------------------------- J 5.0 12 2500 775 50 17.9 17.0 16.4 14.4 H 5.0 12 2500 1775 50 17.8 16.9 16.3 14.3 K 5.0 12 2500 13460 50 16.3 15.4 14.8 12.8 K' 5.0 12 2500 4890 50 16.5 15.6 15.0 13.0 Kn 5.0 12 2500 4890 50 16.4 15.5 14.9 12.9 [Fe 11] 10.0 6 2500 224 100 15.4 14.5 13.9 11.9 H20 10.0 6 2500 978 100 14.7 13.8 13.2 11.2 Cont2.2 10.0 6 2500 1700 100 15.7 14.8 14.2 12.2 CO 10.0 6 2500 7420 100 14.6 13.7 13.1 11.1 Cont3.28 0.3 200 2500 706500 3 10.0 9.1 8.5 6.5 nbL 0.3 200 2500 706500 3 10.7 9.8 9.2 7.2 (n_frames = 2) each with an integration time T = 180 sec and with the object placed at different positions along the slit in both images. The spectral resolution at 2.2 micrometers is 2.2/2200 = 0.001 micrometers/pixel, or a factor of Lambda = 0.33/0.001 = 330 lower than for Kn. In the spectral direction, each pixel sees signal from an area of sky A_Sky = 1 arcsec x 0.5 arcsec = 0.5 arcsec^2, dispersed by a factor of Lambda more than for Kn, but reduced in flux by Tau_Grism. The background current per pixel (averaged over features in the sky emission spectrum) is therefore: [10^((Z.P.-m_Bkg)/2.5) * A_sky * Tau_Grism * Gain/Lambda = 11.8 e-/sec/pixel Here we have used the observed Kn background brightness, M_Bkg, of 12.4 mag/arcsec^2 , and zero point offset, Z.P., of 20.5 mag, and conversion between electrons and ADU, Gain, of 9 e-/ADU. The noise per pixel is then Noise = (50^2 + 180*(11.8+10))^1/2 = (2500 + 3933)^1/2 = 80 e-/pixel We see from this that shot noise from the dark current makes a significant contribution to the total noise, and that integration times of order 180 sec are required to minimize the read noise contribution. Integration times significantly longer than this are not recommended because sky intensity variation make accurate sky subtraction increasingly difficult, and significant numbers of hot pixels saturate as the integration time is increased further. Normally, spectra are recorded as a nodded pair to allow sky subtraction, and preferably as an ABBA sequence. In the case where a single sky frame is subtracted from an object frame to perform the sky subtraction, the total noise per pixel is increased by (2)^1/2 because the sky frame has the same noise per pixel as the object frame. For our example, the final noise becomes 113 e-/pixel. We define a 5*sigma detection per pixel by requiring a signal-to-noise ratio of 5 per pixel after averaging n_y pixels along the slit and averaging the n_frames object frames. The average object signal per pixel in the dispersion direction is then Signal = 5 * Noise / (n_y)^1/2 /(n_frames)^1/2 = 200 e-/pixel The equivalent Kn imaging signal would be Kn Signal = 200 / Gain / Tau_Grism / Tau_Slit * Lambda * n_y^2 = 470000 ADU which corresponds to a Kn brightness of Kn = 20.5 - 2.5 * log (470000/180) = 12.0 mag From this we predict that a signal-to-noise ratio of 5:1 can be achieved with the HK grism on a Kn 12.0 mag object in the continuum in 6 min of on-source integration. The 5*sigma line detection sensitivity in the same time can be estimated by assuming the line occupies two pixels in the dispersion direction. The line flux is then 3.84 * 10^-14 * 10^(-12.0/2.5) * 0.001 * 2 = 1.2 * 10^-21 W cm^-2 Filters ======= Acquisition System The IMB acquisition system contains either the ICCD acquisition TV, SBIG autoguider, or tip-tilt sensor mounted on an X-Y stage at the bottom of the IMB. These systems view the telescope field in optical light through the dichroic reflexer, and allow offset guiding or selection of bright tip-tilt reference stars. The travel of the X-Y stage is ~ +/-31.2mm in each direction in the f/18 (5 arcsec/mm) telescope beam. A TV filter/focusser unit mounted on the X-Y stage directly in front of the TV camera converts the f/18 telescope beam to f/5.76, provides a pupil mask that acts as an optical sky baffle, allows for focussing of the optical image independent of the infrared image, and contains B, V, R, I, and RG630 filters as well as a clear position. F WHEEL CONTENTS The contents of the five wheels are listed below. Note that as of Oct 1994, it is necessary to set the Utility Wheel to increasing numbers in order to ensure that it is latched correctly. The Lens Wheel should also be set to increasing numbers. These details will be fixed as soon as convenient. Table 12: Aperture Wheel Contents Position Keyword Content ---------------------------------------------------------------- 1 Blank Blank 2 Lslit1 1.0 arcsec x 128 arcsec slit 3 Disk5 5.0 arcsec occulting disk 4 Lslit1.5 1.5 arcsec x 128 arcsec slit 5 Lslit2 2.0 arcsec x 128 arcsec slit 6 SlowClr 0.25 arcsec/pixel baffle 7 Lslit5 5.0 arcsec x 128 arcsec slit 8 Lslitl0 10.0 arcsec x 128 arcsec slit 9 Polar Three 20 arcsec x 120 arcsec slits (for polarimetry) 10 Sslit10 10.0 arcsec x 15 arcsec slit 11 Sslit5 5.0 arcsec x 15 arcsec slit 12 FastClr 0.5 arcsec/pixel baffle 13 Sslit2 2.0 arcsec x 15 arcsec slit, 14 Sslit1.5 1.5 arcsec x 15 arcsec slit 15 Disk2 2.0 arcsec occulting disk 16 Sslit1 1.0 arcsecx 15 arcsec slit Table 13: Upper Filter Wheel Contents Position Keyword Filter lambda_c delta lambda (micrometer) (micrometer) -------------------------------------------------------------------------- 1 Blank Blank ... ... 2 Clear Clear ... ... 3 Helium He I 10830 1.082 0.011 4 PGamma H I P_gamma 1.093 0.010 5 PBeta H I P_beta 1.282 0.015 6 FeII MSO [Fe II] + PK50 1.647 0.018 7 AAOFeII AAO [Fe II] 1.650 0.015 8 H2O H2O 1.996 0.050 9 H2_1_0 H2 1-0 S(1) 2.120 0.027 10 BrGamma H I Br_gamma 2.170 0.022 11 Cont2.2 Continuum 2.210 0.094 12 H2_2_1 H2 2-1 S(1) 2.249 0.024 13 CO CO (delta v = 2) 2.343 0.088 14 Cont1.6 Continuum 1.580 0.012 15 Grid Wire Grid Analyser ... ... 16 Mirror IR85 ... ... Table 14: Utility Wheel Contents Position Keyword Content -------------------------------------------------------- 1 Blank Blank 2 Clear Direct imaging cold stop 3 IJ-grism IJ cross-dispersed grism 4 Blank4 Blank 5 Blank5 Blank 6 H-grism H grism 7 SlowCam Slow camera lens (0.25 arcsec/pixel) 8 JH-grism JH cross-dispersed grism 9 Blank9 Blank 10 Blankl0 Blank 11 J-grism J grism 12 Mask Coronograph pupil mask 13 Blank13 Blank 14 HK-grism HK cross-dispersed grism 15 Wollaston Wollaston polarimeter analyser 16 K-grism K grism Table 15: Lower Filter Wheel Contents Position Keyword Filter lambda_c delta lambda (micrometer) (micrometer) -------------------------------------------------------------------------- 1 Blank Blank ... ... 2 Clear Clear ... ... 3 J J 1.275 0.282 4 H H 1.672 0.274 5 KP K' + PK50 2.124 0.337 6 KN Kn + PK50 2.165 0.330 7 K K 2.224 0.394 8 L L' 3.821 0.602 9 Ice 3.08 micrometer H2O ice 3.077 0.102 10 Dust3.28 3.28 micrometer Dust emission 3.299 0.074 11 Dust3.4 3.4 micrometer Dust emission 3.415 0.072 12 Cont3.6 3.6 micrometer Continuum 3.592 0.078 13 Cont4.0 4.0 micrometer Continuum 3.990 0.052 14 BrAlpha H I Br_alpha 4.051 0.054 15 M M 4.777 0.650 16 PK50 2 mm PK50 ... ... Table 16: Lens Wheel Contents Position Keyword Content --------------------------------------------------- 1 Blank Blank 2 FastCam Fast Camera Lens (0.5 arcsec/pixel) 3 Blank3 Blank 4 Clear Clear (Note K Filter Curve Functions can be found in the CASPIR User's Guide) " END_OBJECT = INSTRUMENT_INFORMATION OBJECT = INSTRUMENT_REFERENCE_INFO REFERENCE_KEY_ID = "MCGREGORETAL1996" END_OBJECT = INSTRUMENT_REFERENCE_INFO END_OBJECT = INSTRUMENT END