The WFPC2 CCDs are thick, front-side illuminated devices, with a format of 800x800, 15x15µm multi-pinned phase (MPP). MPP allows CCD exposure with the total inversion of all phases. The Si-SiO₂ interface, at the surface of the CCD, is pinned at the substrate potential, directing signal charge away from the Si-SiO₂ interface states towards the buried n-channel. Figure 4.1 shows a schematic which illustrates the principle of MPP (modified from Janesick, et al. 1989). The front-side Si-SiO₂ interface significantly affects the performance of CCDs, so MPP operation yields many practical benefits including reduced dark noise, better charge transfer efficiency (CTE), rapid removal of residual images, excellent pixel-to-pixel uniformity, and improved radiation hardness. MPP technology has been demonstrated and characterized in both Loral (Janesick, et al., 1989) and Tektronix devices (Woodgate, et al., 1989). The CCD sensors for WFPC2 were made by Loral in 1991 and processed and packaged for flight at JPL.

The Loral CCDs are illuminated from the 'front' surface, i.e., the light passes through the polysilicon gate structure overlying the 10µm thick active silicon layer. Because the WFPC2 devices are front-side illuminated and supported by a bulk silicon substrate, the CCD surface is flat, which has reduced the uncertainties in the astrometric calibration to about the 1/10 pixel level.

In this section the performance of the WFPC2 CCDs is reviewed, and compared to the earlier WF/PC-1 devices. A summary of device characteristics is given in Table 4.1.

***Table 4.1: Comparison of WF/PC-1 and WFPC2 CCDs.***
Parameter	WF/PC-1¹	WFPC2
Device	TI	Loral
Architecture	Thinned	Thick
Illumination	back-side	front-side
Format	800×800	800×800
Pixel size	152 µm	152µm
UV Phosphor	Coronene	Lumogen
Dark rate	0.03 e^- pixel^-1 s^-1(-87°C)	~0.0045 e^- pixel^-1 s^-1(-88°C)
Read noise	13e^- RMS	5e^- RMS
Full well depth	40,000 e^-	~90,000 e^-
Gain	8e^- DN^-1	7e^- DN^-1 or 14e^- DN^-1
ADC range	12 bits (4096 DN)	12 bits (4096 DN)
Full range (e^-)	~30,000e^-	~53,000e^-
QE 6000Å	50%	35%
QE 2500Å	12%	15%
WFC resolution	0.10" pixel^-1	0.0996" pixel^-1
PC resolution	0.043" pixel^-1	0.0455" pixel^-1

¹WF/PC-1 data are available through the STScI data archive.

4.2 Quantum Efficiency

The Loral CCDs are thick, front-side illuminated devices. This lowers their intrinsic QE, due to the absorption of incident light by the polysilicon electrode structure on the front-side surface of the CCD. We also note that due to its MPP operation, the QE of the Loral devices is stable without maintenance such as UV flooding.

The front surfaces of the CCDs are overcoated with a Lumogen phosphor, which serves as the primary detection medium for photons shortward of about 4800Å, down-converting these to 5100Å - 5800Å. Its long wavelength cutoff (4800Å) is also well matched to a CCD's intrinsic sensitivity. The QE of the four flight WFPC2 CCDs is shown in Figure 4.2, which demonstrates the uniform UV response of 10-15% and a peak optical QE of 40%.

This phosphor coating also produces an enhancement of DQE at visual wavelengths, since it acts as an anti-reflection coating.

4.3 Dynamic Range

Linear full well capacity for these devices, clocked appropriately for the MPP mode, is approximately 90,000e^- pixel^-1. Flight qualified ADCs with higher dynamic range (>12 bits) were not available, so WFPC2 operates the two available ADCs at different gain factors, to take partial advantage of both the low read noise and large available full well depth.

One channel has a gain of 14e^- DN^-1, which significantly undersamples the CCD read noise (5 e^- pixel^-1 RMS), and gives a digital full well of about 53,000e^-. The other channel has a gain of 7e^- DN^-1 which is comparable to the CCD read noise, and saturates at about 27,000e^-. The choice of gain factor is determined by the scientific objective. The 7 e^- DN^-1 channel is best suited for faint object and UV imaging, where the lower CCD read noise will be most effective. For example, it should be used for UV imaging of planets or narrowband imaging of high redshift galaxies. The 14 e^- DN^-1 channel has slightly higher effective read noise due to the quantization granularity, but can be used for programs where a signal level in excess of 27,000e^- is required. Even when imaging faint sources, it may be desirable to retain the high signal-to-noise information on brighter field stars as a PSF reference.

Use of the 14 e^- DN^-1channel also allows reasonable recovery of counts for isolated, saturated point sources by summing over the saturated pixels (assuming that the charge bleeding does not extend to the edges of the CCD). See Gilliland (1994).

4.4 Bright Object Artifacts

4.4.1 Blooming

Blooming up and down a CCD column occurs when more than about 90,000e^- (the full well capacity) are collected in any pixel. When the pixel is full, the charge will flow into the next pixels along the column, and so on. The orientation of the bloomed column(s) on the sky depends on the readout direction of the particular CCD (see Figure 1.1 or Figure 3.12 on page 76) and the roll angle of the spacecraft. This effect is visible in Figure 4.3 which shows a logarithmic stretch of the image resulting from a 100s exposure on a star of V magnitude 2.6 through filter F502N in the PC.

Extreme overexposure of the Loral CCDs is not believed to cause any permanent effects, and therefore the WFPC2 does not have a bright object limit.

The WFPC2 CCDs can be operated in a non-standard mode during the integration phase of an exposure, in order to limit the blooming to only those columns containing the bright sources. This is accomplished by operating the serial transfer register clocks during the integration (using the optional parameter CLOCKS as specified in the Proposal Instructions). See Section 2.6 "Serial Clocks" for details.

4.4.2 Horizontal Smearing

During readout of a badly overexposed image, there is spurious charge detected by the readout electronics. The apparent brightness of the stellar halo is higher to the right of the saturated columns. This is particularly obvious at the bottom of the image in Figure 4.3 which is a region in the shadow of the pyramid edge.

The horizontal "smearing" seen in highly saturated images can be modeled as an exponential function which decays over a few rows after a saturated pixel is encountered. The effect itself temporarily saturates after about ten saturated pixels (subsequent saturated pixels have no effect). The effect is twice as bad with gain 7 e^- DN^-1 than with gain 14 e^- DN^-1. This model only works on very highly saturated stellar images.

In Figure 4.3, the image to the right side of the saturated columns is brighter than the left side; and the brightness increases as the number of saturated columns increases. This effect appears to be a signal which starts at a saturated pixel and decays over the next few rows, wrapping around as it does so. The signal is additive with each successive saturated pixel. Jumps are obvious when the number of saturated columns changes. The problem is a known characteristic of the amplifier electronics, and an effort was made to minimize it during design. The increase in signal in rows with saturated pixels is also seen in the over-scan region (the over-scans are provided in ".x0d" files from the pipeline).

An approach to calibrating out the horizontal smearing is described here. An exponential function fits the effect reasonably well. An appropriate algorithm creates an array to contain the signal model. It searches through the uncalibrated image (with the over-scan region included) in the sequence in which the pixels are read out. When it encounters a saturated pixel, it adds an exponential function to the model array, beginning at that pixel. The function has the form s(x)=Ae^-^x^/^h, where x is the offset from the saturated pixel and only positive x values are included. The half-width, h, and amplitude, A, appear to vary from frame to frame and must be determined on the image itself. As more saturated columns are encountered in a row, the signal intensity builds up in the model image. The image can then be "improved" by subtracting the model from the raw image.

The amplitude and half-width parameters can be obtained by trial and error. The typical parameters vary slightly for each chip. The amplitude per saturated pixel is typically 1.75 DN (gain 7) or 0.2 DN (gain 14). On the other hand the half-width at a gain of 14 is larger (h=1800) than at 7 (h=350). So the total integrated effect is about twice as bad at gain 7. A straightforward application of the above algorithm cleaned up most of the signal in rows which had a few saturated columns, but over-subtracted in rows with a large number. The algorithm can be modified to saturate by making the parameter A, which gives the peak contribution from a single saturated pixel, depend on the current level of the effect: A=A₀*(1-C/C_max). This implies that the correction is never larger than C_max no matter how many saturated pixels are encountered. C_max is approximately 14 DN for a gain of 7 and 10 DN for a gain of 14.

The algorithm gives improvement only on highly saturated stellar images (where the star is saturated to 3 or 4 columns at the edges of the chip). On less saturated data, it over-subtracts significantly. This indicates that the problem is nonlinear, and therefore a general algorithm applicable to all data will be difficult to develop.

4.4.3 Diffraction Effects and Ghost Images

Several other artifacts that are common in saturated stellar images are also obvious in Figure 4.3. The spider diffraction spikes caused by both the OTA spiders and internal WFPC2 spiders are at 45° to the CCD columns in all cameras.

The halo around the stellar image is well above the diffraction limit in intensity. Also there are ghost images which result from internal reflections in the filters and in the field-flatteners. These topics are discussed fully in the next Chapter.

4.5 Residual Image

Residual images are seen in front-side-illuminated CCDs, and are associated with the front-side Si-SiO₂ surface interface. When the full well is exceeded, electrons can become trapped at the Si-SiO₂ interface. This trapped charge is slowly released giving rise to residual images. Inverted phase operation (MPP) allows holes to recombine with the trapped electrons at the front-side interface, and so residual images dissipate in a matter of minutes.

A second potential source of residual images, which occurs only in front-side-illuminated CCDs, is known as residual bulk image (RBI). Long wavelength photons can penetrate deeply enough to produce charge in the substrate. Most of this charge recombines rapidly (due to short carrier lifetimes), but some may diffuse into the epitaxial layer, where it can become trapped in epitaxial interface states. Residual images can occur as this charge is slowly released during an exposure. RBI is temperature sensitive since the bulk trapping time constants decrease with increasing temperature. The WFPC2 CCDs do exhibit RBI, but at -70°C trapped charge rapidly escapes so that residual images disappear within 1000s (currently the CCDs are operated at -88°C). Driven by the WFPC2 electronics, the CCDs recover quickly from large over-exposures (100 times full well or more), showing no measurable residual images a half hour after the overexposure.

For images exposed below the saturation level there is a very weak residual image due to charge trapping and charge transfer efficiency (CTE) problem. Measurements on 1800s dark frames interleaved with 2800s exposures of a star field yield a residual flux of 0.3% ± 0.1% of the original star flux, for stars with fluxes from 65 to 17,000 total counts. For typical star fields observed by WFPC2, these residual images are likely to be a problem only for stars that were saturated in a previous image, or for programs where long exposures in low throughput filters are taken immediately after highly exposed images. Hence, repeated exposures at the same CCD position should not lead to any appreciable systematic offset in photometry. CTE is further discussed in Section 4.12.

4.6 Quantum Efficiency Hysteresis

The problem of quantum efficiency hysteresis (QEH) due to back-side charge accumulation has been reviewed in detail by Griffiths, et al. (1989), and Janesick and Elliot (1991). QEH is not present in the Loral CCDs, because they are front-side illuminated and incorporate MPP operation. This has been verified in component tests at JPL. The absence of QEH means that the devices do not need to be UV-flooded, and so decontamination procedures are planned without the constraint of maintaining the UV-flood.

4.7 Flat Field Response

The flat field response is uniform within a few percent, with the exception of a manufacturing pattern defect which generates a 3% reduction in QE once every 34 rows. This pattern defect is caused by a manufacturing error in producing the CCDs; there was a 0.5µm overlap between adjacent 1024x0.5µm raster scans during the construction of the masks used to fabricate the chips. It is identical in all CCDs. The net effect is that every 34^th row on the CCD is approximately 3% too narrow. Photometry of point sources imaged onto these defects will be affected, since the error conserves counts, while flat fields (which are designed to produce a uniform image from a uniformly illuminated target) will effectively multiply the counts in these rows by 1.03. In applications requiring precision photometry across a wide field, it may be useful to correct the images for this flat field effect before performing photometry. There is also an astrometric offset of approximately 3% of the pixel height (0.003" in the WFCs) every 34 rows. Anderson and King (1999) present a nice discussion of these effects.

WFPC2 flat fields also include instrumental effects such as vignetting and shadowing by dust particles, and illumination variations related to optical geometric distortion. For further discussion see Section 5.11.

The WFPC2 CCDs have an intrinsically uniform flat field response since they are not thinned, so there are no large-scale chip non-uniformities resulting from the thinning process. MPP operation also improves pixel-to-pixel uniformity because charge transfer is driven deep into the buried n-channel, away from the influence of Si-SiO₂ interface states. The WFPC2 CCD flat fields show an overall pixel-to-pixel response having <2% non-uniformity. Figure 4.4 shows a portion of a WFPC2 CCD flat field obtained during quantum efficiency measurements at JPL. The image illustrates the excellent pixel-to-pixel uniformity of the Loral devices. The 34 row defect is clearly visible, and its amplitude of 3% serves to calibrate the gray scale.

4.8 Dark Backgrounds

Low dark noise is one of the benefits of MPP, since inverted phase operation suppresses the dominant source of CCD dark noise production (Si-SiO₂ surface states). The remaining source of dark noise, thermal generation in the silicon bulk, is determined by the quality of the silicon used in chip fabrication. The intrinsic dark rate of WFPC2 CCDs is <0.01 e^- pixel^-1 s^-1 at temperatures below -80°C.

The temperature set-points for the WFPC2 TEC coolers are: -88, -83, -77, -70, -50, -40, -30 and -20 °C. The corresponding approximate median dark rates are given in Table 4.2. For instrument health and safety reasons, GOs cannot command temperature changes.

4.8.1 Sources of Dark Current

The dark current appears to have two components: one from electronic sources in the CCD, and a second component whose strength correlates with the cosmic ray flux. The electronic dark current is ~0.001 e^- s^-1, consistent with the Thermal Vacuum Test data.

The second component of dark current appears only on-orbit, its strength drops towards the edges of each CCD, and it is both chip- and time-dependent. At the current operating temperature, this non-electronic component constitutes up to 80% of the total signal measured in the PC. The fraction and overall level are lower in the other chips, and lowest in WF2. This second component ranges from 0.001 e^-s^-1 (WF2) to 0.005 e^- s^-1 (PC). The edge drop off is shown in Figure 4.5, where the average of lines 200-600 for each chip (with hot pixels rejected) is plotted in e^- s^-1 as a function of column number. The drop near the edge is consistent with luminescence from the CCD windows, shadowed by a field stop mask just in front of the CCD.

***Table 4.2: Dark Count Rates.***
CCD Temperature (°C)	Dark count rate (e^- s^-1 pixel^-1)
-20	10.0
-30	3.0
-40	1.0
-50	0.3
-70	0.03
-77	0.016
-83	0.008
-88	0.0045

A further indication of the possible origin of this second component is the correlation between its amplitude and the cosmic ray activity in the same exposure, as shown in Figure 4.6. For example, the cosmic ray flux in the PC varies from 7x10⁵ to 13x10⁵DN per 1000s, while the total dark signal in the PC varies concurrently between 0.0007 and 0.0010 DN s^-1. Similar, though slightly smaller effects are seen in the WFC CCDs. These clues point to cosmic-ray induced scintillation of the MgF₂ field-flattening windows as a likely source of the second dark current component. This might be caused by impurities in the MgF₂ windows; if so, the window of WF2 must contain substantially less impurities. However, other explanations cannot be completely ruled out at this point.

For the great majority of WFPC2 observations, this effect is negligible. In fact, it is noticeable mainly because the true dark rate is very low at the -88°C operating temperature. However, observations for which the dark current is an important limiting factor, either due to noise or background flatness, will require special handling to remove the signature of the dark current properly, as its amplitude depends on the time-variable cosmic ray flux.

4.8.2 Darktime

As of this writing, the "DARKTIME" keyword in the WFPC2 image headers does not reflect correctly the actual time during which the CCD collects dark current. Instead, DARKTIME is merely set equal to EXPTIME (the exposure time) in the data headers, and this value is used for calibration. The error is small, and usually unimportant, but could be significant for programs aimed at measuring the absolute level of the sky background. The actual darktime in seconds is given by

where t is the requested exposure time in seconds, and n is the number of the CCD (PC1=1, WF2=2, etc.), and int() indicates the next lower integer. A duration of 16.4s is required to clear the CCDs before the exposure begins, and 13.6s is needed to read each CCD after the exposure. External exposures of 180s or longer made with the serial clocks off (CLOCKS=NO; the default setting) suffer an additional 60s of darktime (restart=1). This delay is associated with restarting the serial clocks for readout in exposures where the spacecraft AP-17 processor provides shutter control with loss-of-lock checking. Exposures made with the serial clocks on (CLOCKS=YES) avoid this extra 60s (restart=0).

We note that bias frames contain approximately

seconds of dark current. No attempt is made to subtract this from the bias images when creating calibration files for use in the calibration pipeline. This effect is unimportant for most observations, but could be significant if one averaged many undithered deep exposures of the same field, or if one is interested in measuring the absolute level of the sky background. If the dark current were constant in time, this could be corrected by merely changing the value of DARKTIME used during calibration. However, the hotpixels vary on monthly timescales, so this simple correction is only partially successful.

The timing of dark calibration frames is slightly different from that of external science exposures. Dark calibration frames always have restart=0 in Equation 4.1.

The dark calibration reference file in the pipeline is revised weekly to track variations in the hot pixels. The current method of generating these files is to combine the bright hot pixels from typically five on-orbit dark frames taken over the space of about one week, with the low-level dark current from the average of 120 on-orbit dark frames spanning a much longer time period. This method gives an optimal combination of low noise and accurate tracking of hot pixels. Care is also taken that the same super-bias reference files is used for both science data and generation of the dark reference file, as this tends to reduce the noise in long exposures. (Early dark reference files used a much simpler method, and were typically combinations of about ten dark frames taken over two weeks.)

4.8.3 Dark Current Evolution

The dark current in WFPC2 has had an interesting evolution over the lifetime of the instrument. Figure 4.7 shows the median dark current for the central 400 x 400 pixels of each CCD at gain 7, each taken just after WFPC2's monthly decontamination. Each data point represents the median of five raw 1800s dark frames (after rejection of cosmic rays and bias subtraction, normalized to units of DN/1000sec). As such, this plot reflects the uniform, low-level dark current near the center of each detector. During the first six years the dark current increased approximately linearly with time; the dark current increased by a factor of about 2 in the WFC CCDs and by a factor of ~1.3 in the PC. But after 1998 (MJD > 51200) the dark current leveled-off, and perhaps decreased somewhat.

As mentioned before, there are two primary sources of dark current -- a dominant component which is strongly correlated with the cosmic ray flux in the image (probably due to scintillation in the MgF₂ CCD windows; see Figure 4.6), and a smaller thermal dark current in the CCD itself. The dark current increase seen during early years was smaller in the optically vignetted regions near the CCD edges, which suggests much of this increase is related to scintillation effects in the CCD windows. Moreover, the ratio between the dark current at the CCD edge and the CCD center has remained nearly constant throughout the mission (within a range of ~5%; see WFPC2 ISR 2001-05), even though the dark current itself doubled in WF2, WF3, and WF4. Hence, it seems an inescapable conclusion that most of the long-term evolution is related to scintillation effects and variations in cosmic ray flux.

Long-term changes in the cosmic ray flux are perhaps most easily attributed to the solar cycle. The leveling-off of the dark current ~1998 is coincident with the approaching solar maximum which has the effect of reducing the cosmic ray flux at HST's low Earth orbit. Ground-based cosmic ray detectors show a gradual flux increase from 1992 to 1998, followed by a sharper decrease through early 2004. It is possible that other effects might also play some role. For example, portions of the HST orbit near the South Atlantic Anomaly experience higher cosmic ray rates, and it is possible that changes in the HST scheduling system could produce long-term changes in cosmic ray flux and hence dark current. It is also conceivable that long-term changes in the instrument itself might indirectly influence the sensitivity to scintillation effects (e.g. long-term radiation damage might modify the luminescence of camera components).

The thermal dark current of the CCD may also undergo long-term change (i.e. from radiation damage, etc.), and contribute some minor variation. A small increase in the CCD cold junction temperature was seen early in the mission; however, the temperature change can account for only a very small portion of the increase in dark current.

Since the dark current is generally a minor contributor to the total noise in WFPC2 images, its long-term variation is unlikely to impact the quality of WFPC2 observations, except perhaps in special cases (faint sources observed through narrow-band or UV filters, especially in AREA mode).

We note that the variation in dark signal reported here affects all pixels, and thus is distinct from hot pixels which vary in a more cyclic fashion. The hot pixels are highly localized, and are almost certainly due to radiation-damaged sites on the CCD detectors. Their number and intensity increase continuously, but are significantly reduced during decontamination procedures where the CCDs are warmed to +22°C to clear the CCD windows of contaminants. These "decontaminations" were conducted monthly until June 2003, after which their frequency was reduced to 49-day intervals. Apparently the decontaminations anneal defects in the CCDs which produce hot pixels (see Section 4.11).

4.9 Cosmic Rays

HST is subjected to cosmic rays and protons from the Earth's radiation belts. The cosmic ray signature in the Loral CCDs is essentially the same as was seen in the WF/PC-1 devices. Electron-hole pairs generated in the thicker substrate by cosmic rays (and infrared photons) are usually removed by recombination in the low resistivity substrate material, because electrons do not diffuse efficiently up to the collecting phase.

Cosmic ray events usually deposit significant quantities of charge in more than one pixel. This is due partly to the finite thickness of the CCD detectors, and partly to the less than perfect collection efficiency of each pixel. Figure 4.8 shows a histogram of the number of affected pixels for each cosmic ray event. For the purposes of the figure, a cosmic ray is defined as having a peak pixel value more than 10 DN above the background; and an affected pixel is an attached pixel with a value more than 2 DN above the background. Cosmic ray events do have a clear lower cutoff at around 500 electrons of the total signal.

Cosmic ray events impact scientific imaging with WFPC2 differently from WF/PC-1, the previous generation camera. Firstly, the WFPC2 CCDs have an epitaxial thickness of about 10µm compared to 8µm for the thinned WF/PC-1 device, and a recombination length of 8-10µm in the substrate. These facts lead to a higher total number of electrons being deposited per event. WFPC2 CCDs also have lower read noise, and so the number of cosmic ray events apparently differs from that of the WF/PC-1 CCDs, since low amplitude events are detected. In practice, this means that the number of pixels apparently contaminated by cosmic rays in an image is higher in WFPC2, although the underlying event rate is similar to that experienced in WF/PC-1.

Secondly, stellar images are undersampled and much more difficult to separate from cosmic rays, as is shown in Figure 4.9. Faint stellar images and low energy cosmic rays are often indistinguishable. Long science observations are therefore usually broken into at least two exposures (CR-SPLIT) to ensure that events can be identified.

The average properties of on-orbit cosmic ray events have been determined from examination of several dark exposures, each 2000s long. After bias and dark subtraction, "cosmic rays" were identified in each input frame by first looking for pixels more than 5

above the background, and then including in each event all adjacent pixels more than 2

above the background. Very occasionally, two or more physically separate events will be merged into one as a result of this procedure; visual inspection confirms that in the vast majority of cases, this procedure correctly identifies each event and the area affected by it. The typical value of

was 5 to 6 electrons, including both read and dark noise. The region near the borders of each CCD was excluded in order to avoid edge effects, but all results given here are rescaled to the full area of the CCD.

One difficulty in this measurement is caused by hot pixels, for some of which the dark current has significant fluctuations from frame to frame; these can be mistakenly identified as cosmic rays when the dark current is at a maximum. Single-pixel events constitute 10% of the total number of events identified by our procedure, but at least half of them recur in the same position in several frames, thus identifying them as damaged (hot) pixels, rather than true cosmic rays. Also, unlike the majority of cosmic ray events, single-pixel events tend to have a very small total signal; the majority have a total signal of less than 200 electrons, as expected from hot pixels, while the signal distribution of multiple-pixel events peaks around 1000 electrons. For this reason, single-pixel events have been classified as "bad pixels'' rather than "cosmic rays''. While we cannot exclude that some true single-pixel events do occur, they are very rare, probably less than 2% of the total.

Cosmic ray events are frequent, occurring at an average rate of 1.8 events chip^-1 s^-1. The distribution of the total signal is shown in Figure 4.10; it has a well-defined maximum at about 1000 electrons, and a cut-off at about 500 electrons. The latter is well above the detection threshold used for the above measurements (25 electrons in the central pixel of the cosmic ray), and is therefore undoubtedly real.

The histogram in Figure 4.10 shows the distribution of the total energy of all cosmic ray events. One encouraging feature is the very small number of events below about 30 DN. This low energy drop is well above the energy level of excluded single-pixel events.

A good approximation to the cumulative distribution of events as a function of the total signal is given by a Weibull function with exponent 0.25. This function has the form:

where N is the total number of events which generate a total signal larger than S. The best fit to the observed events gives N₀=1.4 events chip^-1 s^-1, S₀=700 electrons, and

. The fit fails below S₀, and should not be extrapolated to low-signal events. The rate of events with the total signal below 700 electrons is 0.4 events chip^-1 s^-1(i.e. total events per CCD per second is N₀+0.4

1.8).

The number of pixels affected by cosmic ray events in a given exposure is a slightly more sensitive function of the threshold used. While there is a clear drop at low signal for both total and peak signal, neighboring pixels can be affected at low levels. Each event (defined as before) affects an average of 6.7 pixels, for about 12 affected pixels chip^-1 s^-1. For a 2000s exposure, this results in about 24,000 affected pixels, or 3.8% of all pixels. As cosmic rays are expected to be randomly placed, a pair of such exposures would have about 900 pixels affected in both exposures; cosmic ray correction is impossible for such pixels. For a pair of 1000s exposures, about 220 pixels will be affected in both frames.

Cosmic ray activity varies as a function of time, geomagnetic latitude of the spacecraft, and other factors. The average numbers given here are subject to change in individual exposures. After studying about one month's worth of dark exposures, we estimate a total range of about a factor of two in cosmic ray rates.

4.10 SAA and Scheduling System Issues

Changes in the WFPC2 observation scheduling system were introduced early in 1999 primarily in order to increase the scheduling efficiency of HST observations starting with Cycle 8.

First, the South Atlantic Anomaly (SAA) contours used to limit WFPC2 observations were modified slightly. The SAA is a region where irregularities in the Earth's magnetic field cause very high cosmic ray rates. WFPC2 imaging is generally not scheduled near the SAA, so as to avoid excessive cosmic ray hits which degrade images by obliterating data in numerous pixels. These adverse effects are usually minimized by operating each instrument only when HST is outside a designated "SAA avoidance contour." (WFPC2 observations of time-critical phenomena can be taken inside the SAA avoidance contour, if necessary.) Biretta and Baggett (1998) analyzed available WFPC2 data, together with data from Air Force satellites flying in similar orbits, and redefined the WFPC2 SAA avoidance contour. This resulted in a 3% increase in designated SAA-free orbits, which allows better scheduling efficiency, and negatively impacts less than 0.1% of WFPC2 science observations. The current (post-1999) contour is given by the M26 area in Figure 4.11.

Second, WFPC2 visits are limited to a maximum length of 5 orbits. Very long visits (up to an earlier maximum of 8 orbits) have very limited opportunities for scheduling, reduce the efficiency of telescope use, and can cause long delays in execution, with long GO wait times. Shorter visits improve the scheduling opportunities for large proposals. One possible drawback is the lower pointing repeatability across visits; this is significant only for programs with special dithering requirements.

A third change since Cycle 8 is that an extra minute of overhead was added to each orbit, which allows splitting an orbit in the Phase II proposal into two separate spacecraft alignments. This one-minute "efficiency adjustment" allows much more efficient scheduling of HST orbits impacted by the SAA.

4.11 Radiation Damage and Hot Pixels

In low Earth orbit (LEO) the CCDs are subject to radiation damage from the Earth's radiation belts. The WFPC2 CCDs are shielded from energetic electrons and about half the incident energetic protons. Long term radiation damage to the CCDs from high energy protons leads to an increase in dark count rate (mainly from the creation of hot pixels), baseline shifts in the CCD amplifiers, and long term degradation of Charge Transfer Efficiency (CTE). There has not been a significant degradation in the amplifier baselines. CTE is discussed in the Section 4.12. On the other hand, one of the major radiation damage mechanisms is the creation of new Si-SiO₂ interface states, which cause increased dark current rates in affected pixels. In the MPP CCD these states immediately recombine with holes, reducing the gradual increase in dark noise by factors of about 25, compared to normal three-phase CCDs (Woodgate, et al. 1989, Janesick, et al. 1989b).

Figure 4.12 is a histogram of the dark current distribution (in e^- s^-1) for hot pixels. It contains three curves: solid for the histogram of all hot pixels just before a decontamination (April 7, 1995); dashed only for the pixels that were hot just before the decontamination and were not hot at the beginning of the cycle (March 10); and long-dashed for pixels that were hot at the start of the cycle and were fixed by a decontamination. Thus, the dashed line represents the "new" hot pixels, and the long dashed line represents the fixed hot pixels. The fact that these two curves are very similar shows that the number of hot pixels is roughly in equilibrium. The majority of new hot pixels have low dark current. The hot pixels that constitute the accumulated legacy of previous periods, and thus survived one or more decontaminations, include higher-current pixels. The population of hot pixels increases at a rate of approximately 33 pixels CCD^-1 day^-1 above a threshold of 0.02 e^- pixel^-1 s^-1, while the camera remains at the normal operating temperature.

About 80% of the new hot pixels return to a normal state at decontamination events when the CCDs are warmed to a temperature of +22°C for 6-12 hours. There is no evidence that the fraction of hot pixels that returns to normal is related to the length of the decontamination. Of those pixels that are not fixed, about half will be fixed after two or three additional decontaminations. After that, the rate of correction decreases. It is conceivable that all hot pixels will be repaired eventually. At the moment there is no evidence of a significant secular increase in the number of hot pixels, and we have a firm upper limit of 8% on the fraction of hot pixels that are not repaired after several decontamination cycles.

In order to deal with the hot pixel problem, we provide monthly lists of possible hot pixels via the World Wide Web. Look for hot pixels under WFPC2 Instrument News at:
http://www.stsci.edu/instruments/wfpc2/wfpc2_hotpix.html

These lists are best used to flag hot pixels as bad. While we do provide an estimate of dark current for each hot pixel as a function of time, there are indications that the noise in hot pixels is much higher than the normal shot noise, and thus dark current subtraction is unlikely to give good results.

4.12 Photometric Anomalies: CTE and
Long vs. Short

There are two photometric anomalies which have now been extensively characterized. The first effect is due to the imperfect charge transfer efficiency (CTE) of the detectors, which causes sources at high row and column numbers to appear fainter than otherwise because the charge is transferred over a bigger fraction of the chip. This anomaly is increasing with time, especially for faint sources, presumably as a consequence of on-orbit radiation damage. In this section, we provide correction formula which reduce the impact of this anomaly to about 1-3% in typical cases. The second effect, called "long vs. short", causes sources to have a lower count rate - and thus appear fainter - in short exposures than in longer exposures and appears independent of the position on the chip. This nonlinearity is very small (i.e. a few percent) or non-existent for uncrowded fields, with less than ~1000 stars per WFC chip. However, for crowded fields with ~10,000 stars per chip, apparent nonlinearities of tens of percent are possible in extreme cases (e.g., when comparing 10 sec. with 1000 sec. exposures). The most likely explanation is that this effect is the result of an overestimate of the sky measurement in the short exposure due to the presence of scattered light around bright stars. Because the magnitude of the "long vs. short" effect is highly dependent on the parameters of the photometric analysis, no standard correction formula have been provided. Both the CTE and "long vs. short" effect are more fully described below.

4.12.1 Charge Transfer Efficiency

The WFPC2 CCDs have a small but significant charge transfer efficiency (CTE) problem which causes some signal to be lost when charge is transferred down the chip during readout. This has the effect of making objects at higher row numbers (more charge transfers) appear fainter than they would if they were at low row numbers. The effect depends on the temperature of the CCDs. At the original temperature of -76°C, as much as 10-15% of the light within a 0.5" radius aperture around a bright star could be lost for objects at the highest rows. As a result, the CCD operating temperature was changed to -88°C on 23 April, 1994. This reduced the effect to a maximum amplitude of 4% for stars with more than 1,500 total detected electrons. This ~4% amplitude seems to remain in effect for stars up to 20,000 total electrons. However, for fainter stars (few electrons) seen against a low background, the effect appears to have grown much larger (up to tens of percent) over the last 8 years. We also note that the effect depends on the amount of background light on the chip. There is significantly less CTE effect in the presence of even a moderate (several tens of electrons) background. Hence, the effect is not well described by either a constant fractional loss or a constant additive loss per charge transfer, but must be calculated as a function of target counts, background light, and epoch.

Our basic understanding is that CTE problems are caused by electron traps in the CCD's silicon. During the readout process these traps capture charge from the image electron packets as they are clocked across the CCD towards the readout amplifier. After some time delay, the charge is released from the traps, but by that time the affected electron packet has moved away, so the re-deposition occurs at some distance from the electron's original position in the image. Hence this has the effect of producing "tails" on images. We believe that larger electron packets fill a larger volume in the bulk silicon, hence brighter images are able to access larger numbers of traps than faint ones. This simple paradigm also suggests that images with high background levels will tend to have less CTE problems, since the background will fill some of the traps, and prevent them from robbing charge as the CCD is read out.

CTE: photometric effects

The primary observational consequence of CTE loss is that a point source at the top of the chip (Y=800) appears to be fainter than if observed at the bottom of the chip (Y=1), due to loss of electrons as the star is read out down the chip (see Figure 4.13). This is called Y-CTE. There also appears to be a similar, but weaker tendency, for stars on the right side of the chip (X=800) to be fainter (called X-CTE). The effects also depend on the brightness of the star and the background level.

The photometric calibration of the instrument presented in this Handbook is based on Holtzman, et al. (1995b). It has been corrected for CTE by assuming a 4% loss across the 800 rows of the CCD (i.e. 2% correction for CCD centers). All of the frames considered in the primary photometric calibration are short exposures of bright stars. While correction formulae have been developed, as discussed below, the 4% ramp is still a reasonable approximation. Hence, for data taken at -88°C, a 4% correction ramp was applied to the measured 0.5" radius aperture photometry, in the sense that objects at row 800 were made brighter by 4%, but the brightness of objects at the first row was not changed. The correction was applied to bring measurements to the values they would have had in the absence of CTE, or equivalently, the values they would have had if measurements had been made at row 0.

Several studies were done on the photometric effects of the Charge Transfer Efficiency (CTE) problem for WFPC2. This work was based on analysis of observations of the globular cluster

Cen (NGC 5139). The first study provides a set of formulae that can be used to correct for CTE loss when doing aperture photometry, based on a data set taken on June 29, 1996 (Whitmore and Heyer 1997, ISR WFPC2 97-08), reducing the observational scatter in these test data from 4-7% to 2-3%, depending on the filter. The second study found evidence that CTE loss for faint stars has increased with time (Whitmore 1998).

A continuation of this analysis using observations of

Cen confirmed that the CTE loss for WFPC2 was time dependent (Whitmore, Heyer, and Casertano 1999). The datasets cover the time range from April 28, 1994 (shortly after the cooldown to -88°C), to February 1999. For bright stars (i.e., brighter than 200 DN when using gain = 14 e^-/DN; equivalent to 400 DN for gain = 7 e^-/DN) there is only a modest increase in the amount of CTE loss as a function of time. However, for faint stars the CTE loss has increased more rapidly. For example, for very faint stars (i.e., 20-50 DN at a gain of 14 e^-/DN) the CTE loss has increased from 3% to 40% for a star at the top of the chip.

It should be noted that CTE loss is strongly dependent on the background level in an image. Figure 4.14 illustrates CTE losses for background levels ranging from 0.03 to 70 DN/pixel. For example, for faint targets (20 - 50 DN, top left panel) a low background of 0.03 DN/pixel results in ~40% CTE loss at late epochs, while a 14 DN/pixel background produces ~4% loss. The results in the previous Figure 4.13 are based on very short (14s) exposures with very low background. By comparison, a typical WFPC2 exposure (300s in F555W) has ~3 DN/pixel background. Hence, the sky background will significantly reduce CTE loss for most science observations. CTE will primarily affect images in the UV and in narrow band filters, where the background is very low.

An approximate correction for stellar photometry is given by Whitmore, Heyer, and Casertano (1999) as follows for stellar photometry performed with a 2 pixel radius aperture. For

and

they give

Note that these equations are for gain = 7 e-/DN observations, since this is most commonly used for science observations. For gain = 14 e-/DN, multiply CTS_obs and BKG by 2 before using the above equations. For further details, please see Whitmore, Heyer, and Casertano (1999).

Another study analyzed CTE losses and developed formulae to correct them (Dolphin 2000). This paper compares WFPC2 observations with ground based observations of Omega Cen and NGC 2419 and derives CTE corrections using a baseline through March 2000, roughly a year longer than available for a similar study by Whitmore, Heyer, and Casertano (1999, PASP, 111, 1559). In general, Dolphin finds good agreement with the Whitmore et al. results (within a few hundredths of a magnitude) with less scatter in the residuals, except for relatively recent (1998 and later) data at low count levels. Dolphin updated his formula on September 17, 2002. For his most up-to-date formulas, the user is strongly encouraged to check his webpage at http://www.noao.edu/staff/dolphin/wfpc2_calib/. Figure 4.15 shows data on Omega Cen which has been corrected for CTE by the Dolphin (2002) formula. Based on this figure and the comparison presented in Whitmore and Heyer (2002), we find that the Dolphin formula provides better CTE corrections than that of Whitmore et al. formula. Our current recommendation is to use the Dolphin (2002) formula for the CTE loss correction, though caution should always be exercised at the faintest levels (e.g. approximately 4 DN in a 14 sec exposure). We list here Dolphin's "complex" equations which take into account the fact that the magnitude loss per pixel is not constant as the star reads out. These CTE correction equations, expressed in magnitudes of CTE loss, are given as follows:


ct = counts in electrons lct = ln(ct) - 7 bg = sqrt (background²+1) - 10 lbg = ln( sqrt(background²+1) ) - 1 yr = (MJD - 50193) / 365.25 Counts and background are given in electrons.

mag(corr) = mag - X-CTE - Y-CTE

Physical effects of CTE

Late in 1999, efforts were made to better understand the detailed effects of CTE during the read out process (Biretta, Baggett, and Riess 2000). Figure 4.16 illustrates the impact of CTE on a single pixel during the read out process. This image is the average of 700 hot pixels taken from WFPC2 dark frames from late 1999, and it effectively shows the system response to a single bright pixel at the center of a CCD. The CTE problem displaces counts into obvious "tails" extending in both the X and Y directions on the CCD. Three components of CTE can be discerned and characterized by the time delay for trapped charge to be released:

All of these components have the effect of robbing charge from typical small apertures (few pixel radius) used for stellar photometry. (A fourth component of CTE is responsible for long-lived residual images, and will be discussed later.)

The brightness profile along the Y-CTE tail is shown quantitatively in Figure 4.17. While the count levels in the extended tail are low, they still make up approximately 2/3 of the total counts displaced from the hot pixel. Figure 4.17 also illustrates the effect in 1994, and gives a clear indication of the time evolution. Similar measurements made on hot pixels in separate intensity ranges are illustrated in Figure 4.18; the total charge in the Y-CTE tail (in this case for late 1999 and background level ~1 DN) is approximately

where I is the hot pixel intensity in DN at gain 7 e^-/DN. This relationship together with Figure 4.17 and model PSFs can be used to predict stellar CTE, and the results appear to be in fair agreement with observations.

Cosmic rays in images are also impacted by CTE, and provide another useful probe of CTE effects. Much like the hot pixels, CTE causes tails to appear on the cosmic rays. Though cosmic rays themselves have complex shapes, these tails are still manifest as a statistical asymmetry, and this asymmetry can be used as a quantitative measure of CTE (Riess, Biretta, and Casertano 1999).

The total counts in these cosmic ray "tails" is a useful metric of CTE. As shown in the top panel of Figure 4.19, no significant tail is apparent at low Y. But at high Y an exponentially declining tail is readily apparent with an e-fold decay of 2 pixels (indicating that charge is released on the 10's of milliseconds timescale). This Y dependence closely mimics that seen in stellar photometry. These tails are very similar to those seen for hot pixels.

Figure 4.20 displays the temporal dependence of both parallel-read (Y) and serial-read (X) induced-tails for WFPC2 as measured with cosmic rays. This figure shows results from thousands of WFPC2 dark frames, and sharply delineates the degradation of CTE with time. There is even evidence for mild acceleration in the sense that the counts in the CR tails at late epochs are somewhat higher than expected by a linear extrapolation of the early data. The same growth trend is seen in Figure 4.20 for X-CTE tails except the X-tails are much weaker and have presently converged at 1/3 the size of the Y-tails. This is in good agreement with the relative strengths of X to Y stellar CTE measurements (Whitmore, Heyer, & Casertano 1999).We note that using internal data, such as these cosmic rays in dark frames, saves external HST pointed time and provides a better time sampling, compared to more conventional stellar CTE monitoring.

As mentioned above, a fourth component of CTE is manifest as long-lived residual images. These residual images are seen as faint ghost images in exposures following a highly exposed target, and tend to decay with a timescale of roughly 10 to 20 minutes (Biretta and Mutchler 1997; Baggett, Biretta, and Hsu 2000). They usually appear at both the location of the bright target, and in pixels below the target (smaller Y values than target). Figure 4.21 illustrates this phenomenon. The trail below the target is caused by charge which is trapped during read-out of the highly exposed image, which is then slowly released during subsequent exposures. The effect is most pronounced when long exposures in low throughput filters (narrow band or UV filters) immediately follow a highly exposed image (usually a broad band filter). These long-lived residual images may be related to surface traps on the CCD, whereas the other components are more likely related to traps in the bulk silicon.

Investigations reveal that CTE losses to extended sources are not uniform across the source (Riess 2000). Rather, they are proportionally greater on the side of the source which is closer to the read amplifier (i.e., low-Y), decrease in the direction away from the amplifier, and charge is regained at the opposite side (i.e., high-Y) of a source. The portion of an extended source which is far from the amplifier suffers little charge loss because charge traps encountered have been filled and in addition, charge is deferred. Our knowledge of how CTE affects galaxies and other extended sources is still growing and it is difficult at this point to provide a recipe to restore changes to the shape of a source. Nevertheless, we suggest that users consider that the total CTE loss expected for an extended source (Baggett et al. 2001; Whitmore, Heyer, & Casertano 1999) likely applies only to the side of the source near the amplifier (i.e., low-Y), with the opposite side (i.e., high-Y) facing smaller losses.

Mitigating CTE during observations

Observers can use a number of strategies to minimize the effect of CTE loss. Longer individual exposures help by increasing both background and source counts, both of which reduce CTE loss. Users thinking of dithering may wish to take this into account if they are considering shortened exposures to allow for more dither positions.

When observing a target significantly smaller than a single detector, it is advisable to place it towards the bottom of a chip (i.e., near the readout amplifier). For example, the aperture WFALL will place the target near the bottom of Chip 3. (Note, however, that targets larger than about 20" centered on WFALL will be split between chips, which itself may lead to photometric problems.) The resulting data can still be corrected using the CTE correction formulae, and the corrections will be smaller.

For faint point sources on low backgrounds, it is recommended that the target be imaged close to the pyramid apex at pixel location (150,150) to reduce the effects of CTE loss. When placing targets closer to the pyramid apex than this position, one risks the target landing near the vignetted regions and affecting the resulting photometry. For the wide field CCDs, aperture = WFALL is recommended. The aperture reference point for WFALL is at pixel (133,149) on the WF3 chip. Therefore, no movement of the target is required to reduce the effects of CTE loss when using this aperture. For PC1 imaging, it is recommended that a POS TARG be used to move the target from the aperture reference point (420.0,424.5) to the recommended position (150,150) using (POS TARG -12.292,-12.491). Table 4.3 presents recommended POS TARGs to position a target at pixel location (150,150) in the respective CCD chip.

***Table 4.3: Recommended POS TARGs to mitigate CTE.***
Aperture	Reference Point X (pixel)	Reference Point Y (pixel)	POS TARG (arcsec)
PC1-FIX	420.00	424.00	-12.292, -12.491
WF2-FIX	423.00	414.00	27.213, -26.293
WF3-FIX	416.00	424.00	26.536, 27.307
WF4-FIX	423.00	421.00	-27.174, 27.001
WFALL	903.00	904.00	no POS TARG needed

When the very highest possible photometric accuracy is required, another possibility is to include a special calibration observation of

Cen, taken close to the time of the science observations and designed so as to reproduce them as closely as possible in exposure and background levels.

A further possible strategy is to preflash the chip to raise the background level. However, tests indicate that the required level of preflash is so high that in general more is lost than gained by this method (due to overhead times and added noise). A variation of this, called "noiseless" preflash, was tested where a flat field exposure is taken immediately prior to a science exposure.¹ However, it gave only very modest improvements in CTE (Schultz, et al. 2001).

As part of the Cycle 8 through 12 Calibration Plans, we continued to monitor the CTE for point sources by repeating the key observations of

Cen every six months (Proposals 7629, 8447, 8821, 9254, 10076). This will be continued in Cycle 13 (Proposal 10364). We also added observations of a cluster of galaxies (Proposal 8456), which yielded a direct measurement of the effect of CTE for faint extended sources for more typical exposure times and background levels. A proposal in Cycle 10 (Proposal 9255) studied the astrometric effects of CTE.

4.12.2 The Long vs. Short Photometric Anomaly

The so-called "long vs. short" anomaly is a nonlinearity of WFPC2 which causes the recorded count rate to increase with exposure time for a given source - the source thus appears brighter in a long exposure than in a short exposure. Suggestions of this nonlinearity was first discussed by Stetson (1995) and then examined in more detail by Kelson et al. (1996), Saha et al. (1996), and Casertano & Mutchler (1998). More recent studies, however, have disputed the existence of the "long vs. short" problem (e.g., Dolphin 2000). Casertano & Mutchler (1998), Hill et al. (1998), and Dolphin (2000) suggest that the apparent "long vs. short" anomaly may be caused by overestimating the value of the sky by a few electrons in the shorter exposure.

We reexamined the "long vs. short" anomaly (Whitmore & Heyer 2002) using the same dataset as in Casertano & Mutchler (1998), that of the globular cluster NGC 2419 taken in 1997 (Proposal 7619). We have analyzed the stars on two chips, WF2 where the globular cluster was centered and WF4 where the stars are much less densely packed. Note that only chip WF2 was analyzed in Casertano & Mutchler (1998). For the uncrowded field (less than ~1000 stars per chip), the typical values of the "long vs. short" nonlinearity are small (a few percent) or nonexistent for stars in the range of 2-400 DN (i.e. 15-3000 electrons) as measured on the "short" exposure. We do not find the larger values (i.e. tens of percent) predicted by Casertano & Mutchler (1998). A re-analysis of WF2 with its densely-packed star field (~10,000 stars per chip) also indicates, in general, a smaller correction than Casertano & Mutchler. Although very large values of a "long vs. short" anomaly are possible for very faint stars in very crowded fields, it appears that the Casertano & Mutchler formula was tuned to fit the worst cases and, therefore, overestimates the values for more typical cases.

We conclude that the "long vs. short" anomaly is very small or non-existent for relatively uncrowded fields (e.g., less than ~1000 stars per chip). However, we still find evidence for varying levels of a "long vs. short" nonlinearity associated with crowded fields, probably due to the subsequent difficulty of estimating accurately the sky background. In order to investigate this possibility further, we separated the measurement of the local sky and the measurement of the object for the crowded NGC 2419 field (WF2). The top panel of Figure 4.22 shows the ratio of the counts in the object aperture of the 1000 sec. exposure relative to the counts in the 10 sec. exposure, using constant values of the sky equal to 30 DN for the long exposure and 0.30 DN for the short exposure. The dashed lines shows the linear fit to the data, with a slope due to the normal CTE effect. The intercept at Y=0 is within ~2 sigma of the theoretical value of 100. Hence, there appears to be little or no "long vs. short" problem for the ratio of the object observations. However, the ratio of the sky values (bottom panel of Figure 4.22) shows an obvious tendency to fall below the theoretical value of 100. This appears to be the cause of the "long vs. short" anomaly in the crowded NGC 2419 field. The sky values in the 10 sec. exposure appear to be overestimated by about 35%, relative to the predicted values based on the sky measurement of the 1000 sec. exposure. This overestimate of the sky background is enough to produce the large effect observed by Casertano & Mutchler (see Whitmore & Heyer 2002).

Because the sky measurement appears to be the cause of the apparent "long vs. short" anomaly, the magnitude of this effect is highly dependent both on the level of crowding in the image and on the specific reduction parameters (e.g., aperture size) used by the observer. Consequently, we do not provide any set formula to correct this effect. The user is referred to Whitmore & Heyer (2002) for more details on this apparent anomaly and for some possible recommendations for dealing with the most adversely affected datasets.

4.13 Read Noise and Gain Settings

The CCDs and their associated signal chains have readout noise levels (in the absence of signal shot noise or interference) of approximately 5e^-. The analog-to-digital converter is highly accurate, and makes virtually no contribution to the read noise, other than the normal information loss caused by digitization of the signal.

The conversion factors from detected electrons (QE x number of incident photons) to data numbers (DN) are tabulated in Table 4.4, as are read noise and linearity ("gamma" is the power law index relating detected DN to input flux). Note that all calculations of sensitivity in this manual assume gains of 7 and 14 for two gain channels, choices very close to the measured gains. The photometric calibration is based on an assumed exact gain of 14 in all CCDs. The measurements given here were derived from thermal vacuum testing. On-orbit measurements have confirmed that the gain ratios are correct to within a possible systematic error of 1%-which will feed directly into a photometric calibration error for gain 7 data, as most of the photometric calibration was done with gain 14 data. Note that the gain ratios are known much more accurately than the individual gains; they are derived from flat field ratios instead. Also, note that the Phase II proposal instructions refer to the ~14 e^- DN^-1 setting as ATD-GAIN=15.

***Table 4.4: Signal Chain Gains.***
Parameter	Gain	PC1	WF2	WF3	WF4
Noise	"7"	5.2 4 ± 0.30	5.51 ± 0.37	5.22 ± 0.28	5.19 ± 0.36
Noise	"15"	7.02 ± 0.41	7.84 ± 0.46	6.99 ± 0.38	8.32 ± 0.46
Gain	"7"	7.12 ± 0.41	7.12 ± 0.41	6.90 ± 0.32	7.10 ± 0.39
	"15"	13.99 ± 0.63	14.50 ± 0.77	13.95 ± 0.63	13.95 ± 0.70
Gamma	"7"	1.0015 ± 0.0006	1.0015 ± 0.0006	1.0020 ± 0.0006	1.0038 ± 0.0007
	"15"	1.0004 ± 0.0001	1.0023 ± 0.0004	1.0032 ± 0.0006	1.0018 ± 0.0012
14/7	ratio	1.987 ± 0.02	2.003 ± 0.02	2.006 ± 0.02	1.955 ± 0.02

Wide Field and Planetary Camera 2 Instrument Handbook for Cycle 14

Chapter 4:
CCD Performance

4.1 Introduction

4.2 Quantum Efficiency

4.3 Dynamic Range

4.4 Bright Object Artifacts

4.4.1 Blooming

4.4.2 Horizontal Smearing

4.4.3 Diffraction Effects and Ghost Images

4.5 Residual Image

4.6 Quantum Efficiency Hysteresis

4.7 Flat Field Response

4.8 Dark Backgrounds

4.8.1 Sources of Dark Current

4.8.2 Darktime

4.8.3 Dark Current Evolution

4.9 Cosmic Rays

4.10 SAA and Scheduling System Issues

4.11 Radiation Damage and Hot Pixels

4.12 Photometric Anomalies: CTE and
Long vs. Short

4.12.1 Charge Transfer Efficiency

CTE: photometric effects

Physical effects of CTE

Mitigating CTE during observations

4.12.2 The Long vs. Short Photometric Anomaly

4.13 Read Noise and Gain Settings

Wide Field and Planetary Camera 2 Instrument Handbook for Cycle 14

Chapter 4: CCD Performance

4.1 Introduction

4.2 Quantum Efficiency

4.3 Dynamic Range

4.4 Bright Object Artifacts

4.4.1 Blooming

4.4.2 Horizontal Smearing

4.4.3 Diffraction Effects and Ghost Images

4.5 Residual Image

4.6 Quantum Efficiency Hysteresis

4.7 Flat Field Response

4.8 Dark Backgrounds

4.8.1 Sources of Dark Current

4.8.2 Darktime

4.8.3 Dark Current Evolution

4.9 Cosmic Rays

4.10 SAA and Scheduling System Issues

4.11 Radiation Damage and Hot Pixels

4.12 Photometric Anomalies: CTE and Long vs. Short

4.12.1 Charge Transfer Efficiency

CTE: photometric effects

Physical effects of CTE

Mitigating CTE during observations

4.12.2 The Long vs. Short Photometric Anomaly

4.13 Read Noise and Gain Settings

Chapter 4:
CCD Performance

4.12 Photometric Anomalies: CTE and
Long vs. Short