Advanced Camera for Surveys Instrument Handbook for Cycle 14
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7.1 The CCDsChapter 7:
Feasibility and Detector Performance
7.1.1 Detector Properties
7.1.2 CCD Spectral Response
7.1.3 Quantum Efficiency Hysteresis
7.1.4 CCD Long-Wavelength Fringing
7.1.5 Optical Performance
7.1.6 Readout Format
7.1.7 Analog-To-Digital Conversion
7.1.8 Flat Fields
7.2 CCD Operations and Limitations
7.2.1 CCD Saturation: the CCD Full Well
7.2.2 CCD Shutter Effects
7.2.3 Readnoise
7.2.4 Dark Current
7.2.5 Cosmic Rays
7.2.6 Hot Pixels
7.2.7 Charge Transfer Efficiency
7.2.8 UV Light and the HRC CCD
7.3 The SBC MAMA
7.3.1 MAMA Properties
7.3.2 SBC Spectral Response
7.3.3 Optical Performance
7.4 SBC Operations and Limitations
7.4.1 MAMA Overflowing the 16 Bit Buffer
7.4.2 MAMA Darks
7.4.3 SBC Signal-to-Noise Ratio Limitations
7.4.4 SBC Flatfield
7.4.5 SBC Nonlinearity
7.5 SBC Bright-Object Limits
7.5.1 Overview
7.5.2 Observational Limits
7.5.3 How Do You Determine if You Violate a Bright Object Limit?
7.5.4 Policy and Observers' Responsibility in Phase I and Phase II
It is the observers' responsibility to ensure that their observations do not exceed the bright-object count limits stated in Table 7.7.
7.5.5 What To Do If Your Source is Too Bright for Your Chosen Configuration?
7.5.6 Bright-Object Protection for Solar System Observations
ACS employs two fundamentally different types of detectors: CCDs for use from the near UV to the near IR, and a Multi-Anode Microchannel Array detector, known as a MAMA, for use in the ultraviolet. The CCD and the MAMA detectors are used in different ways and impose their own unique limitations on the feasibility of observations performed with them. In this chapter we present the properties of the ACS detectors, describe how to use them to optimize scientific programs, and list the steps you should take to ensure the feasibility of your observations.
7.1 The CCDs
7.1.1 Detector Properties
WFC Properties
The WFC/CCD
consists of two 2048×4096 charge-coupled devices that are sensitive from the near UV to the near IR. They are thinned, backside-illuminated devices manufactured by Scientific Imaging Technologies (SITe). They are butted together along their long dimension to create an effective 4096 × 4096 array with a gap corresponding to approximately 50 pixels between the chips.As with STIS, the CCD camera design incorporates a warm dewar window, designed to prevent buildup of contaminants on the window, which were found to cause a loss of UV throughput for the WFPC2 CCDs. A summary of the
ACS/CCDs' performance is given in Table 7.1. The performance values on read noise and dark current are valid as of June 2004.
HRC
The HRC CCD is a flight-spare STIS 1024×1024 CCD also manufactured by SITe. It is also a thinned, backside-illuminated device, but is coated using a process developed by SITe to provide good quantum efficiency in the near-ultraviolet. The performance characteristics and specifications are given in Table 7.1
7.1.2 CCD Spectral Response
WFC
The spectral response of the unfiltered WFC camera is shown in Figure 4.9. This figure illustrates the excellent quantum efficiency in the visible and near infrared part of the spectrum, along with the violet cutoff imposed by the silver coatings on the optical elements.
HRC
The HRC spectral response is also shown in Figure 4.9. As well as excellent quantum efficiency in the visible and near-infrared part of the spectrum, the sensitivity in the near ultraviolet is improved over that of the STIS CCD by means of the coating.
7.1.3 Quantum Efficiency Hysteresis
Based on current data, the ACS CCDs do not suffer from Quantum Efficiency Hysteresis (QEH)-that is, the CCD responds in the same way to light levels over its whole dynamic range, irrespective of the previous illumination level.
7.1.4 CCD Long-Wavelength Fringing
Like most CCDs, the ACS CCDs exhibit fringing in the red, longward of ~7500Å. The amplitude of the fringes is a strong function of wavelength and spectral resolution.
The fringe pattern can be corrected by rectification with an appropriate flat field. The fringe pattern is a convolution of the contours of constant distance between the front and back surfaces of the CCD and the wavelength of the light on a particular part of the CCD. The fringe pattern has been shown to be very stable in similar devices, as long as the wavelength of light on a particular part of the CCD stays constant. In practice, this means that the fringe pattern is dependent on the spectrum of the light incident on the detector, with the sensitivity to the source spectrum a function of the bandwidth of the filter.
7.1.5 Optical Performance
Testing of the WFC and HRC optics and detectors, following fine alignment activities on-orbit, has shown that the optical quality objectives of the cameras are met. The encircled energy values obtained from observations made in SMOV are given in Table 7.2.
Table 7.2: Encircled energy measurements for the ACS channels
WFC at 632.8nm in 0.25 arcsec diameter HRC at 632.8nm in 0.25 arcsec diameter SBC at 121.6nm in 0.10 arcsec diameter
7.1.6 Readout Format
WFC
Each CCD chip is read out as a 4144 × 2068 array, including physical and virtual overscans. This is made up of 24 columns of physical overscan, 4096 columns of pixel data and then 24 further columns of physical overscan. Each column consists of 2048 rows of pixel data followed by 20 rows of virtual overscan. The orientation of the chip is such that for the grism spectra, the dispersed images have wavelength increasing from left to right in the positive x-direction.
HRC
The HRC chip is read out as a 1062 × 1044 array, including physical and virtual overscans. There are 19 columns of physical overscan, followed by 1024 columns of pixel data and then 19 more columns of physical overscan. Each column consists of 1024 rows of pixel data followed by 20 rows of virtual overscan. As with WFC, the orientation of the chip was chosen so that grism images have wavelength increasing from left to right.
7.1.7 Analog-To-Digital Conversion
Electrons which accumulate in the CCD wells are read out and converted to data numbers (DN) by the analog-to-digital converter (ADC). The ADC output is a 16-bit number, producing a maximum of 65,535 DN in one pixel.
The CCDs are capable of operating at gains of 1, 2, 4 or 8 e-/DN. In principle, use of a lower gain value can increase the dynamic range of faint source observations by reducing the quantization noise; however, in practice this improvement is not significant. Table 7.3 shows the actual gain levels and readout noise in electrons for the 4 WFC amps and the default C amp used for the HRC.
Table 7.3: CCD Gain and Readout Noise
For the WFC, gain factors of 1 and 2 are fully supported, and so are gain values of 2 and 4 for the HRC. The remaining two gain factors for each camera are available but unsupported, i.e. users of the latter modes must plan their own calibration. It is worth noticing that, for the WFC, the readout noise associated with GAIN=2 is on average only 0.28 e- higher per amplifier than that of GAIN=1. The noise increase brought about by the use of GAIN=2 is equivalent to that produced by adding a mere 1.7 e- of noise in quadrature to the noise of the GAIN=1 configuration: when the number of detected photons is larger than 3, the Poisson noise alone on the combination of source and sky would exceed this level. Thus, in terms of readout noise, the advantage of using GAIN=1 is minimal, whereas by adopting the higher gain value one would extend by 0.32 magnitude the ability of doing accurate photometry before saturation, would increase the number of bright unsaturated sources to provide cross-image registration and, for point sources, could perform photometry several magnitudes beyond saturation in some cases. Further information about gain values can be found in
ACS ISR 2002-03andACS ISR 04-01.7.1.8 Flat Fields
WFC
The flat fields for the WFC now combine information from two sources. Ground-based flats were obtained for all filters at S/N of ~300 per pixel. To refine the low-frequency domain of the ground flats, in-flight observations of a rich stellar field with large scale dithers have been analyzed (see
ACS ISR 2002-08). The required L-flat correction is a corner-to-corner gradient of 10-18%, dependent on wavelength. The resulting flat field supports photometry to ~1% over the full WFC field of view.Figure 7.1 shows the corrected WFC ground flats for several broadband filters, note that on the sky a gap of 50 pixels exist between the top and bottom halves that is not shown here. The central donut-like structure is wavelength dependent, where pixels in the central region are less sensitive than surrounding pixels in the blue F435W filter, for example, and more sensitive in the red F850LP filter. For further discussion of WFC flat fields, see ACS ISRs
Figure 7.1: WFC Flat Field2001-11and2002-04.
HRC
As for the WFC, the HRC ground flats were refined using in-flight observations of a rich stellar field with large scale dithers to determine the low-frequency domain of the flat fields. The correction required for the visible filters is a corner-to-corner gradient of 6-12%, dependent on wavelength. For the NUV filters, flats were taken in-flight using observations of the bright earth (see
ACS ISR 2003-02) and include both the pixel-to-pixel and low-frequency structure of the detector response.The current HRC flat fields have S/N ~300 per pixel and support photometry to ~1% over the full HRC field of view. Figure 7.2 shows the corrected HRC ground flats, derived for 6 broadband optical filters. The donut-like structure seen in the WFC response is not found in the HRC flats. For further discussion of HRC flat fields, see ACS ISRs
2001-11and2002-04.7.2 CCD Operations and Limitations
7.2.1 CCD Saturation: the CCD Full Well
The full well capacity for the ACS CCDs is given in Table 7.1 as 84,700 e- for the WFC and 155,000 e- for the HRC. This is somewhat dependent on the position on the chip. If the CCD is over-exposed, blooming will occur. This happens when a pixel becomes full, so excess charge flows into the next pixels along the column. However, extreme overexposure is not believed to cause any long-term damage to the CCDs, so there are no bright object limits for the ACS CCDs. When using GAIN = 2 on the WFC and GAIN = 4 on the HRC, it has been shown that the detector response remains linear to well under 1% up to the point when the central pixel reaches the full well depth. On-orbit tests have demonstrated that when using aperture photometry and summing over the pixels bled into, linearity to
Figure 7.2: HRC Flat Field1% holds even for cases in which the central pixel has received up to 10 times the full well depth (see
ACS ISR 04-01for details).
7.2.2 CCD Shutter Effects
The ACS camera includes a very high-speed shutter, so that even the shortest exposure times are not significantly affected by the finite traversal time of the shutter blades. On-orbit testing reported in
ACS ISR 2003-03, has verified that shutter shading corrections are not necessary to support 1% photometry on either the HRC or WFC. A total of 4 exposure times were found to be in error by up to 4.1%, e.g. the nominal 0.1s HRC exposure is really 0.1041s (updates for reference files have now been made to use the correct values in pipeline processing). No significant differences were found between exposure times controlled by the two shutters (A and B), with the possible exception of non-repeatability up to ~1% on the WFC for exposures in the 0.7 - 2.0 sec range. The HRC provides excellent shutter time repeatability.7.2.3 Readnoise
WFC
We measured the read noise level in the active area and overscan regions for all the amplifiers at the default gain settings. In general the read noise has been constant with time. On June 29, 2003, just after a transit through the South Atlantic Anomaly (SAA) the readnoise in Amp A changed from ~4.9 to ~5.9 e- rms.
Although the telemetry did not show any anomaly in any component of the camera, it is likely that the read noise jump was due to some sort of radiation damage. Amp A is the only amplifier that showed this anomaly. The amplitude of the variation (~1 e-) was the same for GAIN 1 and 2. After the following anneal date the readnoise dropped to ~5.5 and it remained constant for 27 days. After the following two anneal cycles the read noise reached stability at ~5.6 e- rms, approximately 0.7 e- higher than before the change and it has remained constant since then. Figure 7.3 shows the read noise in the image area for amplifier A during the "instability" period. The readnoise of all the other amplifiers have been very stable since launch with post-launch figures almost unchanged from the pre-launch measurements made during ground testing.
Even with a slightly higher read noise in Amp A most of the WFC broadband science observations are sky limited while narrowband observations are primarily read noise limited.
Figure 7.3: Readnoise Jump in WFC Amp A (occurred on June 29, 2003). The vertical dashed lines indicate the annealing dates.
HRC
The read noise is monitored only for the default readout amplifier C at the default gain setting (2 e-/DN). No variations have been observed with time. The read noise measured in the image area (4.80 +/- 0.12) is in agreement with the readnoise figure measured in the two overscan regions and it is comparable to the pre-flight value of 4.74 e-.
7.2.4 Dark Current
The dark current of CCDs operated in a radiative environment is predicted to increase with time. Ground testing of WFC devices, radiated with a cumulative fluence equivalent to 2.5 and 5 years of on-orbit exposure, predicted a linear growth of ~1.5 e-/pix/hr/yr. During the testing the WFC CCDs were operated at ~81 C, about 5 degrees cooler than the operating temperature.
The dark current in ACS CCDs is monitored four days per week with the acquisition of four 1000 seconds dark frames (totaling 16 images per week). Dark frames are used to create reference files for the calibration of scientific images and to track and catalog hot pixels as they evolve. The four daily frames are combined together to remove cosmic rays and to extract hot pixel information for any specific day. The dark reference files are generated by combining two weeks of daily darks in order to reduce the statistical noise. The hot pixel information for a specific day are then added to the combined bi-weekly dark. In order to study the evolution of the dark current with time we calculate the modal dark current value in the cosmic-ray free daily darks. As expected, the dark current increases with time (Figure 7.4). The observed rate of dark current linear growth are 1.4 and 2 e-/pix/hr/yr for WFC1 and WFC2 respectively and 1.8 for the HRC CCD. These rates are in general agreement with the ground testing prediction.
Figure 7.4: Dark rate trend with time for the ACS CCDs
7.2.5 Cosmic Rays
Initial studies have been made of the characteristics of cosmic ray impacts on the two ACS imaging cameras, HRC and WFC. The fraction of pixels affected by cosmic rays varies from 1.5% to 3% during a 1000 second exposure for both cameras, similar to what was seen on WFPC2 and STIS. This number provides the basis for assessing the risk that the target(s) in any set of exposures will be compromised. The affected fraction is the same for the WFC and HRC despite their factor of two difference in pixel areas because the census of affected pixels is dominated by charge diffusion, not direct impacts. Observers seeking rare or serendipitous objects as well as transients may require that every single WFC pixel in at least one exposure among a set of exposures is free from cosmic ray impacts. For the CR fractions of 1.5% to 3% in 1000 sec, a single ~2400 sec orbit must be broken into 4 exposures (4 CR splits of 500 to 600 sec each) to reduce the number of uncleanable pixels to 1 or less. (It is also recommended that users dither these exposures to remove hot pixels as well.)
The flux deposited on the CCD from an individual cosmic ray does not depend on the energy of the cosmic ray but rather the length it travels in the silicon substrate. The electron deposition due to individual cosmic rays has a well defined cut-off with negligible events of less than 500 electrons and a median of ~1000 electrons (see Figure 7.5 and Figure 7.6).
Figure 7.5: Electron deposition by cosmic rays on WFC.
Figure 7.6: Electron deposition of Cosmic Rays on HRC.
The distribution of the number of pixels affected by a single cosmic ray is strongly peaked at 4 to 5 pixels. Although a few events are seen which encompass only one pixel, examination of these events indicate that at least some and maybe all of these sources are actually transient hot pixels or unstable pixels which can appear hot in one exposure (with no charge diffusion) and normal in the next. Such pixels are very rare but do exist. There is a long tail in the direction towards increasing numbers of attached pixels.
Distributions of sizes and anisotropies can be useful for distinguishing cosmic rays from astrophysical sources in a single image. The size distribution for both chips peaks near 0.4 pixels as a standard deviation (or 0.9 pixels as a FWHM). This is much narrower than for a PSF and is thus a useful discriminant between unresolved sources and cosmic rays.
7.2.6 Hot Pixels
The dark current and the "hot" pixels on the ACS CCDs have been studied throughout SMOV and Cycle 11. The hot pixels appear similar to those seen on previous CCDs flown on HST and are likely caused by radiation damage.
The dark current distribution is well described by a Gaussian with a center at 0.0022 e-/sec and rms of 0.0029 e-/sec for the WFC, and 0.0025 e-/sec and rms of 0.0015 e-/sec for the HRC. As expected from experience with earlier HST cameras, very significant tails in these distributions are seen from much "warmer" or "hotter" pixels. We have chosen a conservative limit of 0.04 e-/sec for WFC and 0.08 e-/sec for HRC as a threshold above which we consider a pixel to be "hot" and not part of the normal distribution of pixel dark current. Figure 7.7 and Figure 7.8 show the daily growth of these hot pixels. For WFC we find a growth rate of approximately 1200 new hot pixels per day with dark current greater than 0.04 e-/sec. For HRC the number of new hot pixels per day above the threshold is approximately 90. Because the distribution of dark current in hot pixels is strongly peaked near the threshold, the specific number of such pixels is necessarily a strong function of the chosen threshold. During the HST safing events marked in Table 7.7 and Table 7.8, the ACS thermo-electric coolers (TECs) were not operating. This warming reduced the number of hot pixels, as if an anneal had occurred.
Figure 7.7: Hot Pixel Trends for WFC.
Figure 7.8: Hot Pixel Trends for HRC.
The monthly anneals on HRC heal on average 80-85% of new hot pixels, similar to what is seen with WFPC2 and STIS. The anneals on WFC heal only ~60% of hot pixels, leading to a growing population of permanent hot pixels. About 1% of the WFC FOV is covered by permanent hot pixels each year. While the standard CR-SPLIT approach allows for cosmic-ray subtraction, without additional dithering it will not eliminate hot pixels in post-observation processing. Hence, we recommend that observers who would have otherwise used a simple CR-SPLIT now use some form of dithering instead. For example, a simple ACS-WFC-DITHER-LINE pattern has been developed, based on integer pixel offsets, which shifts the image by 2 pixels in X and 2 in Y along the direction that minimizes the effects of scale variation across the detector. The specific parameter values for this pattern are given in Section 8.4.3 of the
Phase II Proposal Instructions. However, any form of dithering providing a displacement of at least a few pixels can be used to simultaneously remove the effects of cosmic ray hits and hot pixels in post-observation processing.Subtraction of a superdark frame from a science image during pipeline calibration can remove the dark current from hot pixels just as it does for normal pixels. However, hot pixels are often orders of magnitude noisier than normal pixels, which in many cases limits their ability to provide useful measurements of flux. In rare cases (but not without precedents), hot pixels can spontaneously "heal", a circumstance which could create false positive detections in some science programs.
7.2.7 Charge Transfer Efficiency
Charge Transfer Efficiency (CTE) is a measure of how effective the CCD is at moving charge from one pixel location to the next when reading out the chip. A perfect CCD would be able to transfer 100% of the charge as the charge is shunted across the chip and out through the serial register. In practice, small traps in the silicon lattice are able to compromise this process by holding on to electrons, releasing them at a significantly later time (seconds rather than microseconds). For large charge packets (several thousands of electrons), losing a few electrons along the way is not a serious problem, but for smaller (~100 electrons or less) signals, it can have a substantial effect.
CTE is typically measured as a pixel transfer efficiency, and would be 1 for a perfect CCD. The CTE numbers for the ACS CCDs at the time of installation are given in Table 7.4. While the numbers look impressive, remember that reading out the WFC CCD requires 2048 parallel and 2048 serial transfers, so that almost 2% of the charge from a pixel in the corner opposite the readout amplifier is lost.
Table 7.4: Charge Transfer Efficiency measurements for the ACS CCDs at installation time (Fe55 test at 1620 e-)
Also, the CTE numbers are significantly different for images where the pixels have a low intensity compared to those where the intensity is high.
Both the WFPC2 and STIS CCDs have been found to suffer from a significant degradation in CTE since their installation in 1993 and 1997, respectively. More details can be found in the latest versions of the WFPC2 Instrument Handbook and the STIS Instrument Handbook.
At the end of Cycle 11 we performed the first on-orbit calibration of the photometric losses due to imperfect CTE on ACS HRC and WFC. We utilized images of 47 Tucanae from a CTE calibration program to measure the dependence of stellar photometry on the number of parallel and serial transfers. The results are described in Riess et al. (
ACS ISR 2003-09) and are summarized here. For WFC, significant photometric losses are apparent for stars undergoing numerous parallel transfers (y-direction) and are ~1-2% for typical observing parameters rising to ~10% in worst cases (faint stars, low background). The size of the photometric loss appears to have a strong power-law dependence on the stellar flux, as seen for other CCD's flown on HST.The dependence on background is significant but there is little advantage to increasing the background intentionally (e.g., by post-flashing) due to the added shot noise. No losses are apparent for WFC due to serial transfer (x-direction). For HRC, significant photometric losses also arise from parallel transfer (~1% for typical observations, ~5% for worst case) but are not seen for serial transfer. Correction formulae are presented in
Figure 7.9: Projected CTE losses (and equivalently, the size of corrections) for example, science applications described in Table 7.5. The precision of measurements is not limited by the size of the loss but rather its uncertainty. As a rule of thumb we suggest that the ultimate limit of precision will be ~25% of the loss after correction.ACS ISR 04-06to correct photometric losses as a function of a source's position, flux, background, time, and aperture size. Figure 7.9 shows the predicted photometric losses for the WFC due to imperfect parallel CTE as a function of time. These curves require extreme extrapolation and should be used for planning purposes only. Four specific science applications are shown as examples: the measure of the faint end of M31's CMD (GO 9453), the measurement of high-redshift supernovae (GO 9528), the measurement of any PSF whose brightness is the zeropoint (i.e., 1 e-/sec) and a 20th mag star in a narrow band.
Table 7.5: Example Science Applications and their Assumed Parameters
When observing a single target significantly smaller than a single detector, it is possible to place it near an amplifier to reduce the impact of imperfect CTE. Although we do not believe this is necessary for Cycle 14, it is easy to accomplish by judicious choice of aperture and target position or by utilizing POS TARG commands (however be aware that such large POS TARGs are not advisable because they change the fractional pixel shifts of dither patterns due to the geometric distortion of ACS). An alternative means to achieve the placement of a target near the amplifier is by using some of the subarray apertures. For example, WFC1-512 (target will have 256 transfers in X and Y), WFC1-1K and WFC1-2K place the target near the B amplifier (or WFC2-2K for the C amplifier).
7.2.8 UV Light and the HRC CCD
In the optical, each photon generates a single electron. However, in the near UV, shortward of ~3200Å there is a finite probability of creating more than one electron per UV photon (see Christensen, O., J. App. Phys. 47, 689, 1976). At room temperature the theoretical quantum yield, i.e., the number of electrons generated for a photon of energy E > 3.5 eV (~3500 Å) is Ne=E(eV)/3.65. The HRC CCDs quantum efficiency curve has not been corrected for this effect. The interested reader may wish to see Chapter 6 of the STIS Instrument Handbook for details on Signal-to-Noise treatment.
7.3 The SBC MAMA
7.3.1 MAMA Properties
The ACS MAMA detector is the STIS flight spare STF7 and provides coverage from 1150 to 1700Å. The MAMA detector is a photon-counting device which processes events serially. The ACS MAMA only operates in the accumulate (
Figure 7.10: Design of the SBC MAMAACCUM) mode in which a time-integrated image is produced. Unlike the STIS MAMAs, the ACS does not offer the high-resolution (2048×2048) mode or the time-tagged data acquisition. The primary benefits afforded by the STIS and ACS MAMAs, in comparison with previous HST UV spectroscopic detectors such as those of the GHRS and FOS, are high spatial resolution, two-dimensional imaging over a relatively large field of view, and low background for point sources.
Figure 7.10 illustrates the design of the MAMA which has an opaque CsI photocathode deposited directly on the face of the curved microchannel plate (MCP). Target photons strike the photocathode, liberating single photoelectrons which pass into the microchannel plate (MCP). There they are multiplied to a pulse of ~4×105 electrons. The pulse is recorded by an anode array behind the photocathode and detected by the MAMA electronics which process it, rejecting false pulses and determining the origin of the photon event on the detector.
The field electrode, or repeller wire, repels electrons emitted away from the microchannel plate back into the channels. This provides an increase in quantum efficiency of the detector at the price of an increase in the detector PSF halo. The repeller wire voltage is always on for SBC observations.
Table 7.6: SBC Detector Performance Characteristics
Characteristic SBC MAMA Performance Photocathode CsI Wavelength range 1150-1700Å Pixel format 1024×1024 Pixel size 25×25 µm Plate scale 0.034×0.030 arcseconds/pixel Field of view 34.6 x 30.8 arcseconds Quantum efficiency 19.2% @ 1216Å Dark count 1.2×10-5 counts sec-1 pix-1 Global count-rate linearity limit1 360,000 counts sec-1 Local count-rate linearity limit ~350 counts sec-1 pix-1 Visible light DQE < 1.2×10-9 above 400 nm
1Rate at which counting shows 10% deviation from linearity. These count rates are well above the bright-object screening limits.
7.3.2 SBC Spectral Response
The spectral response of the unfiltered SBC is illustrated in Figure 7.11. The peak photocathode response occurs at Lyman-
Figure 7.11: ACS SBC Detective Quantum Efficiency. Its spectral response is defined by the cutoff of the MgF2 window at 1150Å at short wavelengths, and by the relatively steep decline of the CsI photocathode at long wavelengths. Out-of-band QE at longer wavelengths (>4000Å) is <10-8 yielding excellent solar-blind performance.
7.3.3 Optical Performance
The SBC exhibits low-level extended wings in the detector point-spread function (PSF). Sample MAMA detector PSF profiles are shown in Figure 7.12.
7.4 SBC Operations and Limitations
7.4.1 MAMA Overflowing the 16 Bit Buffer
The MAMA is a photon-counting detector: as each photon is recorded, it is placed into buffer memory. The buffer memory stores values as 16-bit integers; hence the maximum number it can accommodate is 65,535 counts per pixel in a given
ACCUMmode observation. When accumulated counts per pixel exceed this number, the values will wrap. As an example, if you are counting at 25 counts sec-1 pixel-1, you will reach the MAMA "accumulation" limit in ~44 minutes.One can keep accumulated counts per pixel below this value by breaking individual exposures into multiple identical exposures, each of which is short enough that fewer than 65,535 counts are accumulated per pixel. There is no read noise for MAMA observations, so no penalty is paid in lost signal-to-noise ratio when exposures are split. There is only a small overhead for each MAMA exposure (see Section 9.2).
Figure 7.12: MAMA Point Spread Function
7.4.2 MAMA Darks
MAMA detectors have intrinsically low dark currents. Ground test measurements of the ACS MAMA showed count rates in the range of 10-5 to 10-4 counts per pixel per second as the temperature varied from 28 to 35oC degrees. The count rate increased by about 30% for one degree increase in temperature. In-flight measurements, taken weekly throughout June and July 2002, show count rates between 8*10-6 and 10-5. These measurements were taken as soon as the MAMA was turned on and were therefore at the lower end of the temperature range. A 10 hour observation in SMOV, long enough for nominal temperatures to be reached yield a dark current of 1.2 x10-5 counts per second per pixel. Monthly monitoring throughout cycle 11 shows the in-flight dark current to be about 9x10-6 counts per second per pixel. For typical SBC operations, in which the detector is turned on for less than two hours, a dark image collected at lower temperatures is more suitable and will replace the current calibration image. This has a mean dark rate of 10-5 counts per second per pixel.
The ACS MAMA has a broken anode which disables the seven rows 599 to 605. There are three dark spots, smaller than 50 microns at positions (334,977), (578,964) and (960,851) and two bright spots at (55,281) and (645,102) with rates which fluctuate but are always less than 3 counts per second.
An example of the dark current variation across the detector can be seen in Figure 7.13 below.
Figure 7.13: MAMA Dark Image
7.4.3 SBC Signal-to-Noise Ratio Limitations
MAMA detectors are capable of delivering signal-to-noise ratios on the order of 100:1 per resolution element (2×2 pixels) or even higher. Tests in orbit have demonstrated that such high S/N is possible with STIS (Kaiser et al., PASP, 110, 978; Gilliland, STIS ISR 98-16.)
For targets observed at a fixed position on the detector, the signal-to-noise ratio is limited by systematic uncertainties in the small-scale spatial and spectral response of the detector. The MAMA flats show a fixed pattern that is a combination of several effects including beating between the MCP array and the anode pixel array, variations in the charge-cloud structure at the anode, and low-level capacitive cross-coupling between the fine anode elements. Intrinsic pixel-to-pixel variations are of order 6% but are stable to <1%.
7.4.4 SBC Flatfield
Figure 7.14: Mama Flat Field
High S/N SBC flat fields were taken on the ground. In-flight observations of a UV-bright stellar field with large scale dithers will be used to refine the low frequency structure of the SBC flats. The ground flat in Figure 7.14 illustrates several features. The low frequency response is extremely uniform except for a change of response that can be seen in the four image quadrants. The rows 601 to 605, disabled due to the broken anode, are clearly displayed as is the shadow of the repeller wire running vertically near column 577. A regular fixed "tartan" pattern is visible showing the effect of the discrete anodes. For further discussion of SBC flat fields, see
ACS ISR 1999-02.7.4.5 SBC Nonlinearity
Global
The MAMA detector begins to experience nonlinearity (photon impact rate not equal to photon count rate) at global (across the entire detector) count rates of 200,000 counts sec-1. The nonlinearity reaches 10% at 360,000 counts sec-1 and can be corrected for in post-observation data processing at the price of a loss of photometric reliability. Additionally, the MAMA detector plus processing software are not able to count reliably at rates exceeding 285,000 count sec-1. For this reason and to protect the detectors, observations beyond this rate are not allowed (see Section 7.5).
Local
The MAMA detector remains linear to better than 1% up to ~22 counts sec-1 pixel-1. At higher rates, they experience local (at a given pixel) nonlinearity. The nonlinearity effect is image dependent-that is, the nonlinearity observed at a given pixel depends on the photon rate affecting neighboring pixels. This property makes it impossible to correct reliably for the local nonlinearity in post-observation data processing. In addition, MAMA detectors are subject to damage at high local count rates (see Section 7.5).
7.5 SBC Bright-Object Limits
STScI has responsibility to ensure that the MAMA detectors are not damaged through over-illumination. Consequently, we have developed procedures and rules to protect the MAMA. We ask all potential users to share in this responsibility by reading and taking note of the information in this section and designing observing programs which operate in the safe regime for these detectors.
7.5.1 Overview
The SBC detector is subject to catastrophic damage at high global and local count rates and cannot be used to observe sources which exceed the defined safety limits. The potential detector damage mechanisms include over-extraction of charge from the microchannel plates causing permanent reduction of response, and ion feedback from the microchannel plates causing damage to the photocathode and release of gas which can overpressure the tube.
To safeguard the detector, checks of the global (over the whole detector) and local (per pixel) illumination rates are automatically performed in flight for all SBC exposures. The global illumination rate is monitored continuously; if the global rate approaches the level where the detector can be damaged, the high voltage on the detector is automatically turned off. This event can result in the loss of all observations scheduled to be taken with that detector for the remainder of the calendar (~1 week). The peak local illumination rate is measured over the SBC field at the start of each new exposure. If the local rate approaches the damage level, the SBC filter wheel will be used to block the light, since there is no "shutter". Also, all subsequent SBC exposures (in the obset) will be lost until a new filter is requested.
7.5.2 Observational Limits
To ensure the safety of the SBC detector and the robustness of the observing timeline, we have established observational limits on the incident count rates. Observations which exceed the allowed limits will not be scheduled. The allowed limits are given in Table 7.7, which includes separate limits for nonvariable and irregularly-variable sources. The global limits for irregular variable sources are a factor 2.5 more conservative than for sources with predictable fluxes. Predictable variables are treated as nonvariable for this purpose. Examples of sources whose variability is predictable are Cepheids or eclipsing binaries. Irregularly variable sources are, for instance, cataclysmic variables or AGN.
Table 7.7: Absolute SBC Count-Rate Limits for Nonvariable and Variable Objects
Irregularly Variable1
1Applies to the phase when the target is brightest.
Table 7.8: Limiting V-band Magnitudes for SBC observations in various filters
1System made of a main sequence late-type star with an O5V star contributing 20% to the total light in the V band. In the UV, the O5 component dominates and sets the same limiting magnitude for companion types A-M. A one magnitude safety factor has been added, as for the O5V case.
2Star with a flux distribution like AG Peg.
7.5.3 How Do You Determine if You Violate a Bright Object Limit?
As a first step, you can check your source V magnitude and peak flux against the bright-object screening magnitudes in Table 7.8 for your chosen observing configuration. In many cases, your source properties will be much fainter than these limits, and you need not worry further.
However, if you are near these limits (within 1 magnitude or a factor of 2.5 of the flux limits), then you need to carefully consider whether your source will be observable in that configuration. Remember the limits in these tables assume zero extinction. Thus you will want to correct the limits appropriately for your source's reddening and the aperture throughput.
You can use the information presented in Section 6.2 to calculate your peak and global count rates. Perhaps better, you can use the ACS Exposure-Time Calculator to calculate the expected count rate from your source. It has available to it a host of template stellar spectrograms. If you have a spectrum of your source (e.g., from IUE, FOS, or GHRS) you can also input it directly to the calculator. The calculator will evaluate the global and per pixel count rates and will warn you if your exposure exceeds the absolute bright-object limits. We recommend you use the ACS exposure time calculator if you are in any doubt that your exposure may exceed the bright-object MAMA limits.
7.5.4 Policy and Observers' Responsibility in Phase I and Phase II
It is the observers' responsibility to ensure that their observations do not exceed the bright-object count limits stated in Table 7.7.
It is your responsibility to ensure that you have checked your planned observations against the brightness limits prior to proposing for Phase I. If your proposal is accepted and we, or you, subsequently determine (in Phase II), that your source violates the absolute limits, then you will either have to change the target, if allowed, or lose the granted observing time. We encourage you to include a justification in your Phase I proposal if your target is within 1 magnitude of the bright-object limits for your observing configuration. For SBC target-of-opportunity proposals, please provide in your Phase I proposal an explanation of how you will ensure your target can be safely observed.
STScI will screen all ACS observations that use the MAMA detector to ensure that they do not exceed the bright-object limits. In Phase II, you will be required to provide sufficient information to allow screening to be performed.
Here we describe the required information you must provide.
Prism Spectroscopy
To allow screening of your target in Phase II for spectroscopic MAMA observations you must provide the following for your target (i.e., for all sources which will illuminate the detector during your observations):
- V magnitude.
- Expected source flux at observing wavelength.
- Spectral type (one of the types in the screening tables).
- E(B-V).
- B-V color.
If you wish to observe a target which comes within one magnitude (or a factor of 2.5 in flux) of the limits in the spectroscopic bright-object screening table (Table 7.8) for your configuration, after correction for reddening, but which you believe will not exceed the absolute limits in Table 7.7 and so should be observable, you must provide auxiliary information to justify your request. Specifically:
- You must provide an existing UV spectrum (e.g., obtained with IUE, FOS, GHRS, STIS or GALEX) of the star which proves that neither the global nor the local absolute limits will be exceeded.
Imaging
The SBC imaging bright-object screening magnitudes are very stringent, ranging from V = 15 to V = 20.5 for the different imaging apertures, and apply to all sources imaged onto the MAMA detector (i.e., not just the intended target of interest). Table 7.8 can be used to determine if the target of interest is above the bright-object limit. Starting in Cycle 8, STScI has been using the second-generation Guide-Star Catalog (GSC II) to perform imaging screening for objects in the field of view other than the target itself. The GSC II contains measurements from photometrically calibrated photographic plates with color information for magnitudes down to at least V = 22 mag. This information will be used to support bright-object checking for fixed and for moving targets (major planets). STScI will make a best effort to perform the imaging screening using GSC II. However, observers should be prepared for the possibility that under exceptional circumstances GSC II may be insufficient. For instance, fields close to the Galactic plane may be too crowded to obtain reliable photometry. If for any reason the screening cannot be done with GSC II, the observer is responsible for providing the required photometry. In the case of moving targets, STScI will identify "safe" fields, and the observations will be scheduled accordingly. Observers will be updated on the status of their observations by their Program Coordinators. We anticipate that bright-object considerations will not have a significant effect on the scheduling of such observations.
Policy on Observations Which Fail Because they Exceed Bright-Object Limits
If your source passes screening, but causes the automatic flight checking to shutter your exposures or shut down the detector voltage causing the loss of your observing time, then that lost time will not be returned to you; it is the observer's responsibility to ensure that observations do not exceed the bright-object limits.
7.5.5 What To Do If Your Source is Too Bright for Your Chosen Configuration?
If your source is too bright, there may be no way of performing the observation with the SBC. The SBC has no neutral-density filters and only low resolution prism dispersing modes. The options open to you if your source count rate is too high in a given configuration include:
- Change configurations totally to observe a different portion of the spectrum of your target (e.g., switching to the CCD).
- Attempt to locate an equivalent but less bright target.
7.5.6 Bright-Object Protection for Solar System Observations
Observations of planets with ACS require particularly careful planning due to the very stringent overlight limits of the SBC. In principle Table 7.7 and Table 7.8 can be used to determine if a particular observation of a solar-system target exceeds the safety limit. In practice the simplest and most straightforward method of checking the bright object limits for a particular observation is to use the ACS Exposure-Time Calculator. With a user-supplied input spectrum, or assumptions about the spectral energy distribution of the target, the ETC will determine whether a specified observation violates any bright object limits.
Generally speaking, for small (<~0.5-1 arcsec) solar-system objects the local count rate limit is the more restrictive constraint, while for large objects (>~1-2 arcsec) the global limit is much more restrictive.
As a first approximation, small solar system targets can be regarded as point sources with a solar (G2V) spectrum, and if the V magnitude is known, Figure 7.7 and Table 7.8 can be used to estimate whether an observation with a particular ACS prism or filter is near the bright-object limits. V magnitudes for the most common solar-system targets (all planets and satellites, and the principal minor planets) can be found in the Astronomical Almanac. This approximation should provide a conservative estimate, particularly for the local limit, because it is equivalent to assuming that all the flux from the target falls on a single pixel, which is an overestimate, and because the albedos of solar-system objects in the UV are almost always significantly less than their values in the visible part of the spectrum (meaning that the flux of the object will be less than that of the assumed solar spectrum at UV wavelengths where the bright-object limits apply). A very conservative estimate of the global count rate can be obtained by estimating the peak (local) count rate assuming all the flux falls on one pixel, and then multiplying by the number of pixels subtended by the target. If these simple estimates produce numbers near the bright-object limits, more sophisticated estimates may be required to provide assurance that the object is not too bright to observe in a particular configuration.
For large solar-system targets, checking of the bright-object limits is most conveniently done by converting the integrated V magnitude (Vo, which can be found in the Astronomical Almanac) to V magnitude/arcsec2 as follows:
where area is the area of the target in arcsec2. This V / arcsec2 and the diameter of the target in arcsec can then be input into the ETC (choose the Kurucz model G2 V spectrum for the spectral energy distribution) to test whether the bright- object limits can be satisfied.
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