HST Data Handbook for WFPC2
 |
 |
 |
 |
 |
|
| HST Data Handbook for WFPC2 | ||||
|
5.2 Photometric Corrections
A number of corrections must be made to WFPC2 data to obtain the best possible photometry. Some of these, such as the corrections for UV throughput variability, are time dependent, and others, such as the correction for the geometric distortion of WFPC2 optics, are position dependent. Finally, some general corrections, such as the aperture correction, are needed as part of the analysis process. We describe each class in turn.
5.2.1 Time-Dependent Corrections
The most important time-dependent correction is that for the contamination of the CCD windows, which primarily affects UV observations. Other time-dependent corrections are due to the change in operating temperature in April 1994 and the variations of the PSF with focus position; the latter is also position-dependent (see section "Aperture Correction" for more information).
Contamination
Contaminants adhere to the cold CCD windows of the WFPC2. Although these typically have little effect upon the visible and near infrared performance of the cameras, the effect upon the UV is quite dramatic and can reduce throughput for the F160BW filter by about 30% after 30 days. These contaminants are largely removed during periodic warmings (decontaminations) of the camera, and the effect upon photometry is both linear and stable and can be removed using values regularly measured in the WFPC2 calibration program. Table 5.4 shows the monthly contamination rates measured for each detector (
WFPC2 ISR 98-03). Table 5.5 provides decontamination dates up until November 2001; anupdated list of decontamination dates is maintained on the WWW.Contamination is typically measured from the bimonthly observations of the WFPC2 primary standard, the white dwarf GRW+70d5824 (see table 5.4 for rates), and is therefore directly applicable to blue objects. The early observations for the standard photometric filters (F336W, F439W, F555W, F675W, and F814W) have been supplemented by observations of a stellar field in the globular cluster
Cen (mean B-V ~0.7 mag); the contamination rates measured (in parentheses in table 5.4) are generally in good agreement with those measured for GRW+70d5824. The
WFPC2 ISR 96-04). These results will be verified further with the analysis of UV observations of NGC 2100, a young globular cluster in the LMC.Based on standard star observations taken 1994 to 1998, the long-term photometric throughput appears to be quite stable: fluctuations are ~2% or less peak-to-peak in filters longwards of and including F336W (
ISR 98-03; and theWWW monitoring memo). However, the UV throughput has gradually been evolving; specifically, from 1994 to 1998, the clean count rates (measurements immediately after the decons) in some filters, e.g., F160BW by ~12% in the PC, and F170W by ~9%, while decreasing slightly in others, e.g., F255W by ~3%. Furthermore, the monthly contamination rates slowed slightly for the UV filters, e.g., from ~1% per day to 0.5% per day in F160BW on the PC (see table 5.4 andISR 98-03). Both the monthly and long-term contamination rates reported in the ISR have been incorporated into the WFPC2 contamination tables in synphot.There are three ways observers may correct UV data for the variable decline rates and small zeropoint shifts; note that in general, filters redward of F336W do not require a correction based on day since decon but may benefit from a small zeropoint correction. The first option is to use the synphot contamination tables (see example below), which contain both the monthly contamination effects as well as the long term zeropoint changes. Note however, that synphot linearly interpolates in both wavelength and date to obtain the final contamination estimate; the only wavelengths currently in the synphot contamination tables are the central wavelengths of the monitoring filters. The second option is to use the rates in table 5.4 to correct for contamination based upon the number of days since a decon procedure; the
ISR 98-03provides details on correcting for the small zeropoint offset. The third method is to use theWWW table of monitoring resultsto obtain data close in time and wavelength to the science observation and derive an independent model of the contamination effects.Observers should be aware that the calibration pipeline does not automatically correct for contamination. The correction must be applied manually (e.g., using results tabulated in
WFPC2 ISR 98-03) or using synphot. The example below illustrates how to use the calcphot task to determine the effect of contamination in an observation. In this case, the command computes the expected count rate for a WF3, F218W observation taken 20 days (MJD=49835.0) after the April 8, 1995, decontamination, with the gain=7 setup. Removing thecont#49835.0from the command will determine the count rate if no contamination was present. Note that the entire photometric mode must be in double quotes, otherwise STSDAS treats the "#" as a comment indicator and will ignore the remainder of the command line. An 8000 K black body spectrum was chosen largely as a matter of simplicity-the correction values for contamination depend only on the filter chosen and do not reflect the source spectrum.Table 5.4: Contamination Rates (Percentage Loss per Day)
sy> calcphot "wfpc2,3,f218w,a2d15,cont#49835.0" \
>>> spec="bb(8000)" form=counts
Filter Year1 % decline per day error % decline per day error % decline per day error % decline per day error F160BW 4/94-4/95 -0.885 0.104 -- -- -1.277 0.038 -- -- F160BW 4/95-4/96 -0.840 0.137 -1.281 0.089 -1.227 0.057 -0.930 0.102 F160BW 4/96-4/97 -0.690 0.159 -1.303 0.094 -1.142 0.069 -1.079 0.067 F160BW 4/97-6/98 -0.613 0.138 -1.202 0.077 -1.093 0.065 -0.893 0.067 F170W 4/94-4/95 -0.564 0.009 -0.949 0.011 -0.988 0.009 -0.801 0.012 F170W 4/95-4/96 -0.516 0.012 -0.901 0.012 -0.956 0.011 -0.736 0.013 F170W 4/96-4/97 -0.509 0.012 -0.901 0.011 -0.943 0.012 -0.752 0.012 F170W 4/97-6/98 -0.452 0.013 -0.804 0.011 -0.772 0.012 -0.643 0.012 F218W 4/94-4/95 -0.478 0.028 -- -- -0.846 0.019 -- -- F218W 4/95-4/96 -0.433 0.036 -0.757 0.039 -0.787 0.026 -0.704 0.047 F218W 4/96-4/97 -0.442 0.044 -- -- -0.845 0.033 -- -- F218W 4/97-6/98 -0.398 0.046 -- -- -0.633 0.047 -- -- F255W 4/94-4/95 -0.236 0.008 -- -- -0.471 0.007 -- -- F255W 4/95-4/96 -0.235 0.007 -0.445 0.013 -0.490 0.007 -0.385 0.013 F255W 4/96-4/97 -0.215 0.008 -0.398 0.017 -0.333 0.013 -0.298 0.011 F255W 4/97-6/98 -0.183 0.007 -0.311 0.014 -0.327 0.012 -0.234 0.012 F300W 4/95-4/96 -- -- -0.260 0.015 -- -- -- -- F336W 4/94-4/95 -0.031 0.007 -- -- -0.198 0.007 -- -- F336W 4/95-4/96 -0.108 0.008 -0.188 0.013 -0.208 0.009 -0.226 0.014 (-0.127) (0.060) (-0.143) (0.030) (-0.153) (0.027) (-0.157) (0.023) F336W 4/96-4/97 -0.068 0.009 -0.286 0.014 -0.285 0.010 -0.153 0.010 F336W 4/97-6/98 -0.062 0.008 -0.177 0.012 -0.181 0.010 -0.053 0.010 F439W 4/94-4/95 -0.008 0.008 -- -- -0.078 0.008 -- -- F439W 4/95-4/96 -0.037 0.007 -0.121 0.013 -0.080 0.009 -0.127 0.014 (-0.007) (0.047) (-0.073) (0.023) (-0.077) (0.030) (-0.077) (0.023) F439W 4/96-4/97 -0.037 0.007 -0.146 0.017 -0.083 0.011 -0.031 0.011 F439W 4/97-6/98 -0.026 0.005 -0.147 0.017 -0.094 0.012 -0.014 0.014 F555W 4/94-4/95 -0.018 0.007 -- -- -0.067 0.007 -- -- F555W 4/95-4/96 -0.039 0.005 -0.050 0.013 -0.051 0.009 0.016 0.016 (-0.023) (0.043) (-0.023) (0.023) (-0.030) (0.030) (-0.027) (0.027) F555W 4/96-4/97 -0.028 0.004 -0.055 0.011 -0.012 0.009 0.023 0.009 F555W 4/97-6/98 -0.010 0.004 -0.048 0.012 -0.031 0.011 -0.020 0.011 F675W 4/94-4/95 -0.009 0.008 -- -- -0.009 0.008 -- -- F675W 4/95-4/96 -0.012 0.006 -0.009 0.013 0.006 0.009 -0.020 0.016 (-0.067) (0.067) (-0.003) (0.037) (-0.007) (0.037) (-0.013) (0.037) F675W 4/96-4/97 -0.034 0.007 -- -- 0.000 0.000 -- -- F675W 4/97-6/98 0.024 0.006 -- -- -- -- 0.013 0.016 F814W 4/94-4/95 0.041 0.007 -- -- 0.005 0.008 -- -- F814W 4/95-4/96 -0.014 0.006 -0.001 0.012 -0.038 0.009 0.045 0.015 (-0.043) (0.063) (-0.007) (0.030) (-0.000) (0.030) (-0.007) (0.033) F814W 4/96-4/97 -0.005 0.005 -0.035 0.015 -0.005 0.011 0.012 0.010 F814W 4/97-6/98 -0.000 0.003 -0.087 0.012 0.014 0.010 -0.043 0.010
1 Epoch boundaries are at Apr. 25, 1994 (MJD=49467), Apr. 25, 1995 (MJD=49832), Apr. 25, 1996 (MJD=50198), Apr. 25, 1997 (MJD=50563), and Jun. 29, 1998 (MJD=50993).
Table 5.5: Dates of WFPC2 Decontaminations through November 2001.1
1994 1997 2000 1998 2001
Cool Down on April 23,1994
The temperature of the WFPC2 was lowered from -76° C to -88° C on April 23, 1994, in order to minimize the CTE problem. While this change increased the contamination rates, it also improved the photometric throughput, especially in the UV, reduced the CTE losses, and greatly reduced the impact of warm pixels. Table 5.6 provides a partial list of corrections to table 5.1 for the pre-cool down throughput. Including the MJD in a synphot calculation (as shown in appendix A of the
Table 5.6: Ratio Between Pre- and Post-Cool Down ThroughputSynphot User's Guide) using up-to-date synphot contamination tables will correct PHOTFLAM for this change.
PSF Variations
The point spread function (PSF) of the telescope varies with time, and these variations can affect photometry that relies on very small apertures and PSF fitting. Changes in focus are observed on an orbital timescale due to thermal breathing of the telescope and due to desorption, which causes a continual creeping of the focal position. This change had been about 0.7 µm per month until mid-1996, when it greatly slowed. Currently the focus drift is less than 0.3 µm per month. The effect of focus position on aperture photometry is described in
WFPC2 ISR97-01. About twice a year, the focal position of the telescope is moved by several microns to remove the effect of the desorption; thefocus historyis available online.In addition, jitter, or pointing motion, can occasionally alter the effective PSF. The Observatory Monitoring System (OMS) files provide information on telescope jitter during observations (see appendix C). These files are now regularly provided to the observer with the raw data. Observations taken after October 1994 have jitter files in the Archives. Limited requests for OMS files for observations prior to October 1994 can be handled by the STScI Help Desk (E-mail
help@stsci.edu).Remy et al. (1997) have been able to obtain high-quality photometry of well-exposed point sources by modeling the point spread function with TinyTim (Krist, 1995), and taking into account focus and jitter terms via a chi-squared minimization method. Similar results have been obtained using observed PSFs (Surdej et al., 1997), provided that the PSF used is less than 10" from the observed star and corresponds to a spectral energy distribution similar to that of the target. The
WFPC2 PSF librarywas established to help users find suitable PSFs, if they exist, or carry out experiments with what is available.5.2.2 Position-Dependent Corrections
In this Section we discuss the Charge Transfer Efficiency (CTE) correction and the possibly related long vs. short anomaly, the geometric distortion, the gain differences between different chips, and the effect of pixel centering.
Charge Transfer Efficiency
The CTE problem (or CTI, charge transfer inefficiency) results in targets losing counts as the WFPC2 chips are read out. The effect is most pronounced for objects near the top (Y~800) of the chip where the more rows the target must be clocked through, the more charge is lost. The problem has been attributed to impurities in the silicon which trap the charge, preventing it from being read out immediately (Holtzman et al., 1995). During on-orbit calibrations early in WFPC2's mission, the Investigation Definition Team (Holtzman et al., 1995) discovered the presence of CTI in the CCDs and at that time, with an operating temperature of -76oC, it was measured at ~10-15%. Lowering the operating temperature reduces the CTI, so the camera temperatures were set as low as possible, -88o C, which reduced the CTE (charge transfer efficiency) losses to ~3-4% in the worst case (top of the chip, in 1994).
After the acquisition of several more years of calibration data, the CTI was found to be increasing over time (Whitmore & Heyer, 1997). Stellar aperture photometry of
Effect on Point Sources
Since the April 1994 cool down, the CTE loss rate has been steadily worsening. Extensive observations made during Cycles 5 through 9 have provided a monitor of the CTE calibration over time; while the effect was about 5% in 1996-1997, the losses are up to 50% or more (peak-to-peak) in 2001 for the faintest targets2. Recent publications (PASP, 112, 1397, PASP 111, 1559, and WFPC2 ISRs
01-09,97-08) quantify the CTE effect under various observational circumstances and provide empirically-derived correction formulae. After these corrections, the residual CTE effect for well-exposed stars is estimated to be less than 2%.The correction for CTE depends on a number of variables, including the average background, the average counts over the chip, the counts in the source itself, and the target position on the chip. Assuming a 2 pixel radius aperture, the corrected counts are given by:
where X and Y are the coordinates of the star center in pixels, and XCTE and YCTE are the percentile loss over 800 pixels in the x and y direction, respectively. For observed counts less than 4000 DN and backgrounds greater than 0.1 DN, the Y and X CTE components are given by:
For observed counts higher than 4000 DN and blank sky backgrounds greater than 0.1 DN, the Y and X CTE contributions are as follows:
where MJD is the Modified Julian Date and BKG is the mean number of counts in DN for a blank region of the background.
For details on the use of these formulae, see Whitmore et al. (PASP 111, 1559). Note that the equations are for gain = 7, the mode most commonly used for science observations. Multiply CTSobs and BKG by 2 before using either set of equations for gain=15. Also note that due to small uncertainties in the bias level, negative BKG can sometimes occur when the background is very low (i.e.,
0.05 DN). For BKG < 0.1 DN, Whitmore et al. (PASP 111, 1559) recommend linearly scaling the value based on observations using the same filters but with longer exposures of the same field. When longer exposures are not available and the background is blank sky, approximate background levels can be estimated by scaling values provided in Table 1 of the reference given above, though changes in zodiacal light and scattered Earth light may limit their accuracy. Another option is to set BKG to some threshold value (e.g., 0.1 DN) though this will provide only a rough approximation of XCTE and YCTE.
Other correction formulae for CTE losses in point source photometry have been developed by Dolphin (PASP, 112, 1397). This new study compared WFPC2 observations with ground based observations of
Effect on Extended Sources
The effect of CTE on the shape and structure of extended sources has also been studied (
ISR 00-04). Pairs of images of individual galaxies observed near and far from the read-out amplifier were subtracted. The average profile of such galaxy residuals is distinctly asymmetric and indicates that the charge is primarily lost from the leading (amplifier-side) of the galaxy image. The side of the galaxy away from the amplifier suffers little charge loss because the charge traps encountered during the read-out have already been filled (by the leading edge) and because some trapped charge is being released. Preliminary results of aperture photometry on faint galaxy cores show effects consistent with those predicted for CTE in stellar photometry.The analysis of how CTE affects galaxies is still at a relatively early stage and it is difficult at this point to provide a quantitative correction which can be employed by WFPC2 observers. Nevertheless, observers may wish to consider that the total CTE loss expected for a galaxy applies only to the half of the galaxy near the amplifier, with the other half experiencing negligible losses.
Long vs. Short Anomaly (non-linearity)
A comparison of repeated images of the same stellar field can appear to give higher count rates for the faint stars on longer exposures. The effect was first noted by Stetson, as reported in Kelson et al. (1996) and Saha et al. (1996). These authors used a 5 % correction factor for the effect in their papers. A more detailed study was made by Casertano and Mutchler, as reported in
ISR 98-02. They found that this apparent non-linearity appears to be strictly a function of total counts in each source. The magnitude errors produced are less than 1% for well-exposed stars (over 30,000 e-), but may be much larger for faint stars. However, subsequent studies have not been able to confirm the existence of the so called "long vs. short" anomaly (e.g., Saha, Labhart, & Prosser 2000; and Dolphin 2000).Based on a recent study by Whitmore and Heyer, it now appears that the effect may only be present for very crowded fields, where accurate sky subtraction becomes difficult or impossible. Hence, it may be largely a function of the photometric parameters and algorithms that are used for a particular study. On typical fields, where the stars are not so crowded that they overlap, the effect appears to be less than a few percent, if it exists at all. Further investigation of this topic is continuing and the reader is referred to a new ISR on the subject that will be out in early 2002.
Geometric Distortion
Geometric distortion near the edges of the chips results in a change of the surface area covered by each pixel. The flatfielding corrects for this distortion so that surface photometry is unaffected. However, integrated point-source photometry using a fixed aperture will be affected by 1 to 2% near the edges, with a maximum of about 4-5% in the corners. A correction image has been produced and is available from the Archive (
f1k1552bu.r9h). The counts measured for a star centered at a given pixel position must be multiplied by the value of these same pixels in the correction image (simply multiplying your image by the correction image and then measuring the counts is easiest). A small residual effect, due to the fact that the aperture radius differs from the nominal size, depends on the aperture used and is generally well below 1%. Please see section 5.4 for additional discussion on geometric distortion and astrometry.The 34-th Row Defect
Approximately every 34th row on the WFPC2 detectors possesses a reduced sensitivity (Trauger et al., 1993); these features are likely due to a manufacturing defect that resulted in the affected rows being somewhat narrower than the rest. The pipeline flatfields contain these features and produce calibrated images appropriate for surface brightness measurements. Point source photometry, however, can suffer 1-2% errors; and, the narrow rows have a significant effect on astrometry, causing periodic errors of up to 0.03 pixel. A recent paper by Anderson & King (1999) provides photometric and astrometric correction formulae for this 34th row effect.
Gain Variation
The absolute sensitivities of the four chips differ somewhat. Flatfields have been determined using the gain=14 setup, normalized to 1.0 over the region [200:600,200:600]. However, most science observations are taken using the gain=7 setup. Because the gain ratio varies slightly from chip to chip, PHOTFLAM values will be affected. The count ratios for the different chips from Holtzman (1995b) are:
These count ratios should be included in the zeropoint calculation if using values from Holtzman et al. (1995b) on gain=7 data. Conversely, their reciprocals should be applied when using synphot zeropoints on gain=14 data. If you use the value of PHOTFLAM from the header to determine your zeropoint, the different gains for the different chips will already be included. Remember to use the new PHOTFLAM values provided in table 5.1 or the post-May 1997 synphot tables; those included in the header for data taken before May 1997 will have less accurate values.
Pixel Centering
Small, sub-pixel variations in the quantum efficiency of the detector could affect the photometry. The position of star relative to the sub-pixel structure of the chip is estimated to have an effect of less than 1% on the photometry. At present there is no way to correct for this effect.
Possible Variation in Methane Quad Filter Transmission
Based on results from Jupiter and Uranus archival WFPC2 data, the extended wings of the methane filter transmission curve appear to vary across the field of view (Karkoshka, priv. comm.). While this is unimportant for objects with flat spectra, it can have a major impact on photometry of objects with methane bands, where a significant fraction of photons comes from the wings. To provide data to check the methane filter, a set of eight 40-sec Saturn images in a 3x3 grid around the FQCH4W3 methane quad filter aperture (one of the 9 positions falls outside of the filter) is planned as part of the Cycle 10 calibration program (proposal 9256). The magnitude and direction of the effect will be quantified by comparing results from the rings of Saturn (flat spectrum) to results from Saturn itself (deep methane band spectrum).
Anomalous Rotational Offset in the Linear Ramp Filters
For completeness, this effect is included here though we expect no impact on observations as any photometric effect is estimated to be less than 1%. Analysis of FR533N VISFLAT images has revealed an apparently randomly occurring offset of about 0.5 degrees in the filter wheel rotation for some images, a quantity that corresponds to one filter step. The pivot point of the rotation implicates the filter wheel as the source of the inconsistency. A handful of other filters, on different filter wheels as well, appear to exhibit the same problem; at this time, the source of this anomaly, whether it is mechanical or due to a software error, is unknown. A detailed report is in progress (Nov. 2001); a preliminary report is available in
ISR 01-04.5.2.3 Other Photometric Corrections
Miscellaneous corrections that must be taken into account include: aperture corrections, color terms if transforming to non-WFPC2 filters, digitization noise and its impact on the estimate of the sky background, the effect of red leaks and charge traps, and the uncertainty of exposure times on short exposures taken with serial clocks on.
Aperture Correction
It is difficult to measure directly the total magnitude of a point source with the WFPC2 because of the extended wings of the PSF, scattered light, and the small pixel size. One would need to use an aperture far larger than is practical. A more accurate method is to measure the light within a smaller aperture and then apply an offset to determine the total magnitude. Typically, magnitudes will be measured in a small aperture well-suited to the data at hand-a radius of 2-4 pixels, with a background annulus of 10-15 pixels, has been found adequate for data without excessive crowding-and the results corrected to the aperture for which the zeropoint is known. The aperture correction can often be determined from the data themselves, by selecting a few well-exposed, isolated stars. If these are not available, encircled energies and aperture corrections have been tabulated by Holtzman et al. (1995a). If PSF fitting is used, then the aperture correction can be evaluated directly on the PSF profile used for the fitting.
For very small apertures (1-2 pixels), the aperture correction can be influenced by the HST focus position at the time of the observation. The secondary mirror of HST is known to drift secularly towards the primary and to move slightly on time scales of order of an orbit. The secular shift is corrected by biannual moves of the secondary mirror, but the net consequence of this motion is that WFPC2 can be out of focus by up to 3-4 µm of secondary mirror displacement at the time of any given observation. This condition affects the encircled energy at very small radii, and thus the aperture corrections, by up to 10% in flux (for 1 pixel aperture in the PC); see
WFPC2 ISR97-01for more details. If the use of very small apertures is required-because of crowding, S/N requirements, or other reasons-users are strongly advised to determine the aperture correction from suitable stars in their images. If such are not available, an approximate aperture-focus correction can be obtained as described inWFPC2 ISR97-01.A standard aperture radius of 0."5 has been adopted by Holtzman et al. (1995b; note that Holtzman et al. 1995a used a radius of 1."0). For historic consistency, the WFPC2 group at STScI and the synphot tasks in STSDAS refer all measurements to the total flux in a hypothetical infinite aperture. In order to avoid uncertain correction to such apertures, both in calibration and in science data, this infinite aperture is defined by an aperture correction of exactly 0.10 mag with respect to the standard 0."5 aperture. This value (0.10 mag) is close to the values tabulated in Holtzman et al. (1995a) as well as consistent with Whitmore (1995); only extended sources larger than ~1" are likely to require a more accurate aperture correction (
ISR 97-01). Equivalently, the total flux is defined as 1.096 times the flux in the standard aperture of 0."5 radius. In practice, this means that observers wishing to use our tables or the synphot zeropoints should:
- Correct the measured flux to a 0."5 radius aperture.
- Apply an additional aperture correction of -0.10 mag (equivalently, multiply the flux by 1.096).
- Determine the magnitude using the zeropoints given.
See also the example in section 5.2.4.
Color Terms
In some cases it may be necessary to transform from the WFPC2 filter set to more conventional filters (e.g., Johnson UBV or Cousins RI) in order to make comparisons with other datasets. The accuracy of these transformations is determined by how closely the WFPC2 filter matches the conventional filter and by how closely the spectral type (e.g., color, metallicity, surface gravity) of the object matches the spectral type of the calibration observations. Accuracies of 1-2% are typical for many cases, but much larger uncertainties are possible for certain filters (e.g., F336W, see Red Leaks section below), and for certain spectral types (e.g., very blue stars). Transformations can be determined by using synphot, or by using the transformation coefficients in Holtzman et al. (1995b).
Digitization Noise
The minimum gain of the WFPC2 CCDs, 7 e-/ADU, is larger than the read noise of the chip. As a result, digitization can be a source of noise in WFPC2 images. This effect is particularly pernicious when attempting to determine sky values, because the measured values tend to cluster about a few integral values (dark subtraction and flatfielding cause the values to differ by slightly non-integral amounts). As a result, using a median filter to remove objects that fall within the background annulus in crowded fields, can cause a substantial systematic error, whose magnitude will depend on the annulus being measured. It is generally safer to use the mean, though care must then be taken to remove objects in the background annulus.
A more subtle effect is that some statistics programs assume Gaussian noise characteristics when computing properties such as the median and mode. Quantized noise can have surprising effects on these programs. Based upon the analysis of a variety of possible strategies for sky determination (
WFPC2 ISR 96-03), the centroid and the optimal filter ("ofilter") of the histogram of sky pixel values were found to produce the least biased result for typical WFPC2 data with low background levels.Red Leaks
Several of the UV filters have substantial red leaks that can affect the photometry. For example, the U filter (F336W) has a transmission at 7500 Å that is only about a factor of 100 less than at the peak transmission at about 3500 Å. The increased sensitivity of the CCDs in the red, coupled with the fact that most sources are brighter in the red, makes this an important problem in many cases. The synphot tasks can be used to estimate this effect for any given source spectrum.
Charge Traps
There are about 30 macroscopic charge transfer traps, where as little as 20% of the electrons are transferred during each time step during the readout. These defects result in bad pixels, or in the worst cases, bad columns and should not be confused with microscopic charge traps which are believed to be the cause of the CTE problem. The traps result in dark tails just above the bad pixel, and bright tails for objects farther above the bad pixel that get clocked out through the defect during the readout. The tails can cause large errors in photometric and astrometric measurements. In a random field, about 1 out of 100 stars are likely to be affected. Using a program which interpolates over bad pixels or columns (e.g., wfixup or fixpix) to make a cosmetically better image can result in very large (e.g., tenths of magnitude) errors in the photometry in these rare cases. See also section 3.5.4.
Exposure Times: Serial Clocks
The serial clocks option (i.e., the optional parameter CLOCKS = YES in the Phase II proposal instructions) is occasionally useful when an extremely bright star is in the field of view, in order to minimize the effects of bleeding. However, the shutter open time can have errors of up to 0.25 second due to the manner in which the shutters are opened when CLOCKS=YES is specified. If the keyword SERIALS = ON is in the image header, then the serial clocks were employed. Header information can be used to correct this error. The error in the exposure time depends on the SHUTTER keyword. If the value of this keyword is "A", then the true exposure time is 0.125 second less than that given in the header. If instead the value is "B", then the true exposure time is 0.25 second less than the header value.
Users should also note that exposure times of non-integral lengths in seconds cannot be performed with the serial clocks on. Therefore, if a non-integral exposure time is specified in the proposal, it will be rounded to the nearest second. The header keywords will properly reflect this rounding, although the actual exposure time will still be short as discussed above.
F1042M Extended Halo
Observers using the F1042M filter should be aware that it possesses an anomalous PSF containing additional light in a broad halo component. The defocused halo is likely due to the CCD detector becoming transparent at this wavelength, so that the light is reflected and scattered by the back of the CCD. The scattering will impact photometry in the F1042M filter relative to other filters, since a greater fraction of the counts will lie outside the 1 arcsecond diameter aperture used for photometry on standard stars. The
WFPC2 Instrument Handbookprovides a comparison of azimuthal averages for an observed F1042M and F953N PSF; additional PSF examples are in theWFPC2 PSF Library.5.2.4 An Example of Photometry with WFPC2
This example shows the steps involved in measuring the magnitude of the star #1461 (Harris et al., 1993) in the Cousins I passband. The image used for this example can be obtained from the HST Archive, or from the
Table 5.7: Magnitude of Star #1461 inWWW. The WWW directory contains the materials forWFPC2 ISR95-04, A Demonstration Analysis Script for Performing Aperture Photometry. Table 5.7 shows the results from an analysis script similar toWFPC2 ISR95-04, but including some of the corrections discussed above.
Images: u2g40o09t.c0h[1] and u2g40o0at.c0h[1]
Position: (315.37,191.16)
Filter: F814W
Exposure Time: 14 seconds
Date of observation: MJD - 49763.4
2113.49 counts Raw counts in 0.5" radius aperture (11 pixels for PC) -48.63 = 2064.86 counts Background subtraction (0.12779 counts x 380.522 pix obtained from a 40-pixel radius aperture with an annulus of 5 pixels) x 0.9915 = 2047.31 counts Correction for geometric distortion. Not needed if doing surface photometry. => 15.481 mag Raw magnitude (=-2.5 x log10(2047.31 / 14 sec) + 20.894) NOTE: -2.5 x log10(1.987) has been added to the zeropoint from section 5.2.2, Gain Variation since these calibrations were taken using the gain=14 setup. Most science observations use gain=7. -0.10 = 15.381 mag Aperture correction to total magnitude, estimated from Holtzman (1995a). -0.028 = 15.353 mag CTE correction (using formulas 1, 2d, 3d from WFPC2 ISR97-08with this data). Note that this example uses relatively old CTE corrections which do not include a time dependence. Please see section 5.2.2 for the most up-to-date corrections.-0.000 => mF814W = 15.353 mag Contamination correction (0.000 x [49763.4 - 49760.1]). An additional correction to transform to e.g., Cousin I, could be applied using the method described in Contamination.
1 Alist of decontamination dates is kept updated on the WWW.
2 In general, typical WFPC2 exposures are much longer than the short calibration images, resulting in a higher background which significantly reduces the CTE loss and minimizes the CTE problem for most science exposures.
|
Space Telescope Science Institute http://www.stsci.edu Voice: (410) 338-1082 help@stsci.edu |
| |
|
|
|
|
