The QE of the flight IR detector, as measured at the Goddard Detector Characterization Lab (DCL), is shown as a solid curve in Figure 5.19
. The QE curve demonstrates very high sensitivity of the IR detector for wavelengths longer than 1000 nm. The actual total system throughput of WFC3 depends on many factors including the HST
OTA, pick off mirror, filter transmission functions, QE, etc. Based on ground measurements of these quantities, the total system throughput was calculated and compared to the first on-orbit measurements. A 5–20% increase in the total system throughput was discovered, which we attribute to multiple factors. The dashed curve represents the QE under the assumption that the entire flight correction is in the QE. Note, however, that this assumption is unphysical given the realities of anti-reflection coatings and interpixel capacitance.
WFC3 IR exposures taken with an aluminum blank in place, rather than a filter, provide a measure of the detector dark current. The dark current of the flight array has a skewed distribution, with a mode, median, and mean of 0.045, 0.048, and 0.048 e–
/s/pixel respectively. The shifted mode is due to the asymmetry of the dark-current distribution among the pixels, characterized by a long tail of “hot pixels” randomly located across the detector. The mean dark current has remained unchanged in the first three years of in-flight operations (WFC3 ISR 2012-11
The histogram of dark current values, along with the cumulative dark-current distribution, i.e., the fraction of pixels with a dark current lower than a certain level, is shown in Figure 5.20
. (see WFC3 ISR 2009-21
for further details on dark current calculations). Note that in broad filters, the zodiacal light background is 0.3-1.0 e–
/s/pixel, a factor of 10-20 times larger than the dark current. The WFC3 ETC can be used to compute the zodiacal light contribution for a given pointing, in addition to providing thermal and dark current estimates.
The IR detector has four independent readout amplifiers, each of which reads a 512×512 pixel quadrant. The four amplifiers generate very similar amounts of read noise. This is illustrated in Figure 5.21
(left), which compares the correlated double sampling (CDS) read noise levels for the four quadrants of the detector. CDS read noise refers to the noise associated with subtracting a single pair of reads. These read noise values were derived from a series of RAPID ramps taken during SMOV testing, providing a measure of the total noise in a difference image. For short ramps, such as these RAPID ramps, the contribution of shot noise due to dark current accumulation is less than 0.01 e–
. Figure 5.21
(left) therefore shows that the CDS read noise of the detector is between 20.2–21.4 e–
By averaging over multiple reads, the effective noise of an IR ramp can be reduced. As shown in Figure 5.21
(right plot), the effective noise in a SPARS200 ramp can be reduced from ~20.0 e–
down to ~12.0 e–
(2 reads plus zeroth read and 15 reads plus zeroth read, respectively). Similar reductions in noise can be achieved with other sample sequences (WFC3 ISR 2009-23
Before launch, ground-based flats were obtained for the 15 imaging IR filters at a mean S/N of ~500 per pixel using an external optical stimulus (WFC3 ISR 2008-28
). Because the overall illumination pattern of the ground-based flats did not precisely match the illumination attained on-orbit from the OTA, there are errors in these ground-based flats on large spatial scales. These errors were initially measured by performing stellar photometry on rich stellar fields that were observed using large-scale dither patterns during SMOV and cycle 17. In the SMOV exposures for 4 of the wide (W) filters, the rms difference between the sigma-clipped average magnitude of a star and its magnitude in the first pointing was 1.5%, independent of wavelength (WFC ISR 2009-39
). The errors have since been determined more accurately by creating sky flats from thousands of on-orbit exposures, masking out astronomical sources. Flat field reference files corrected using these sky flats were delivered in December 2011. These reference files are expected to support photometry to better than 1% rms accuracy over the full WFC3 IR field of view. A detailed description of their production and accuracy is given in WFC3 ISR 2011-11
. Analysis of a grid of observations of a standard star in 3 filters has shown that these flats produce consistent photometry over most of the detector, contributing rms uncertainty ~0.007 mag to photometric measurements (WFC3 ISR 2013-01
shows examples of bias-corrected ground-based flats taken with wide-band filters (left: F110W, right: F160W). Both flats are displayed with an inverse greyscale stretch chosen to highlight features.
The WFC3 IR calibration program shows that the detector response is in fact (slightly) non-linear over the full dynamic range. This behavior is illustrated in Figure 5.23
, which presents a plot of average counts as a function of time. The black diamonds are the measured average signal; a linear fit has been made to the signals up to ~25,000 electrons (solid red line). The dashed red line shows this best-fit line extended out to the total exposure time of the ramp. The blue horizontal line marks the level at which the counts deviate by more than 5% from linearity (about 78,000 electrons). For the purposes of non-linearity correction, the 5% nonlinearity level has been defined as “saturation.”
Previous HgCdTe detectors on HST
have suffered from a count-rate dependent non-linearity. We are investigating this effect on the WFC3-IR detector. An initial measurement of this effect was made by comparing the photometry of star clusters observed over a wide dynamic range and at overlapping wavelengths in WFC3-IR and NICMOS and/or ACS-WFC. We found a significant detection of a non-linearity in WFC3-IR photometry which is in the same direction but a few times smaller than that of NICMOS. From the stars we measured a non-linearity of WFC3-IR of ~1% per dex over a range of 10 magnitudes (4 dex) which was independent of wavelength. (See WFC3 ISR 2010-07
.) This measurement was confirmed using exposures that boosted count rates with Earth limb light (WFC3 ISR 2010-15
) and observations of groups of stars observed with 2MASS (WFC3 ISR 2011-15
). The impact of this non-linearity is that photometry of faint (i.e., sky dominated) sources calibrated with WFC3-IR zeropoints will appear 0.04 +/-0.01 mag too faint.
The make-up of the WFC3/IR detector’s pixel population includes several flavors of anomalously responsive pixels: hot, cold, unstable, dead, and deviant in the zeroth read. Hot pixels, those showing excess charge, are defined as pixels with more than 100 times the average dark current. Cold pixels are inversely sensitive to incident photons and exhibit a negative slope when measured up the ramp (i.e., pixel value is lower in last frame up the ramp compared to first frame). The anomalous response of a cold pixel could be due to lower intrinsic QE in that pixel or to surface defects. Unstable pixels, as the name implies, are those that behave in an unpredictable fashion; that is, the signal up the ramp does not repeat reliably from ramp to ramp (see Appendix 2, WFC3 ISR 2010-13
for examples). There are dead, or unbonded, pixels which do not respond to light (Figure 5.24
). Overlapping the dead pixel population is the population of pixels which have bad zeroth read values, generally due to being short-circuit or unbonded (WFC3 ISR 2003-06
The level of the IR crosstalk is only ~1e−06
that of the target flux (WFC3 ISR 2010-02
); for unsaturated sources, the crosstalk is below the background noise. Once a source saturates, the crosstalk becomes visible at about the level of the background and remains constant as the voltage of the device is pinned.
Several examples of persistence in WFC3 observations and strategies for avoiding persistence are described in Section 7.9.4
. A description of a phenomenological model of persistence used to aid in removing the effects of persistence is given in the WFC3 Data Handbook
shows the characteristic shape of persistence versus fluence as observed in a series of darks following an image of Omega Cen which had been deliberately exposed to a level where many stars in the image were saturated. The first dark exposure took place a few minutes after the end of the Omega Cen exposure and the last dark exposure took place about one orbit later. (See WFC3 ISR 2013-07
.) The amount of persistence is fairly small until the exposure level reaches about half of full well and saturates near full well exposure. The persistence gradually decays with time from the first dark exposure (highest curve in figure) to the last dark exposure (lowest curve in figure).
Persistence decays roughly as a power law with time, as illustrated in Figure 5.26
, which is based on the data displayed in Figure 5.25
. The different curves here show the decay for different levels of saturation, as measured in electrons. Persistence at low fluence levels decays more rapidly than persistence at high fluence levels. There are 3 curves for each level corresponding to the 3 times this experiment was repeated. The differences are partially due to the fact that different pixels were illuminated to different levels each time, but may also indicate some intrinsic variability that is not understood. For comparison, the dark current is about 0.015 electrons/sec. If one assumes that a power law describes persistence from 100 to 10,000 seconds after an exposure, then one concludes that about 3% of charge is trapped in an exposure that has a nominal fluence level of 100,000 electrons.