For broad band filters the sky background will limit the detection of faint targets. For example, an 8-orbit observation in F555W gives a ~5

detection limit at Johnson V=28.6 for an average sky level of 23 mag arcsec^-2in V. Note that the sky background is a strong function of position, especially for targets near the ecliptic; the sky level can vary from V=23.3 mag arcsec^-2 at the ecliptic pole to about V=20.9 mag arcsec^-2 on the ecliptic near the solar avoidance limit. (See Table 6.3 on page 160 for sky level as function of ecliptic coordinates.)

If these higher sky levels would severely impact the science data, observers should consider specifying the special requirement LOW-SKY on the Phase II proposal. This parameter forces the observation to be made when the sky background is within 30% of the minimum value for the target. Note, however, that this will also reduce the number of HST calendar windows available to the observation, and so could result in scheduling delays or may even make the observation infeasible if there are other constraints such as ORIENTs. A minor decrease in the per-orbit visibility period also results from LOW-SKY, but for background limited programs this is a minor price to pay for the guarantee of a much lower background. In summary, LOW-SKY will reduce the sky background, but should only be used if the science goals require it.


	*Note that LOW-SKY cannot be used for CVZ targets, as they imply mutually exclusive pointing constraints.*

Scattering of bright Earth light in the OTA can produce non-uniformities in the background which may hamper analysis of faint target images. Most often these take the form of diagonal bars of suppressed background light in several of the CCDs. These effects tend to occur for broad band filters when the OTA axis is about 25° from the bright Earth. This effect is most often seen in observations of targets in the CVZ (continuous viewing zone), since the Earth limb is never very far from the OTA axis when observing in the CVZ. Figure 7.1 shows a typical case. LOW-SKY will eliminate this effect for non-CVZ targets, as it places the OTA axis more than 40° from the bright Earth. Alternatively, one can place the target away from the CCD center to avoid these artifacts.

Another option for reducing the sky brightness, is the special requirement SHADOW, which forces the observation to be made when HST is in the Earth's shadow. This usually has a large negative impact on the observing efficiency, and is recommended only when observing far-UV emission lines (e.g. Ly

and OI 1304Å). Its primary goal is only to reduce geocoronal emission lines. Moreover, it does not attempt to minimize zodiacal emission, which dominates at visible wavelengths.

7.2 Observing Bright Targets

Saturation is the primary concern when observing bright targets. The analog-to-digital converter will run out of bit codes at ~28,000 e^-pixel^-1 in the ATD-GAIN=7 (default) setting, and at ~53,000 e^-pixel^-1 in the ATD-GAIN=15 setting. Count levels above these are merely reported as 4095 DN in the data. Hence ATD-GAIN=15 is recommended for targets approaching 28,000 e^-pixel^-1. The disadvantage of this setting is that the read noise is poorly sampled by this coarse digitization, and hence the read noise is slightly increased.

At count levels above ~90,000 e^-pixel^-1 charge will overflow the potential well of each pixel, and begin to bloom up and down the CCD columns. For example, this occurs in the F555W filter at about V=13.5 for a 10s exposure on the WFC, and at about V=13.0 on the PC1.

At very high count levels, above ~10⁸ e^-per CCD column, the charge bloom will reach the top and bottom of the CCD and flow into the serial registers. CLOCKS=YES will dispose of this charge as it reaches the ends of the CCD, and thus prevent it from leaking back into adjacent CCD columns. This exposure level corresponds roughly to a 10s exposure of a V=7 star in F555W. Note that CLOCKS=YES offers no benefit unless the bloom reaches the ends of the CCD, and that it may slightly compromise the bias and dark calibration. Moreover, CLOCKS=YES will result in anomalous exposure times; exposure times are rounded to the nearest integral second, minus a delay time of up to 0.25s for the shutter to open. (See Serial Clocks for further discussions about the use of CLOCKS=YES.)

Besides setting ATD-GAIN=15, the PC CCD can be used to reduce saturation effects for stellar objects. The peak of the PSF will be spread over more pixels on the PC (vs. WFC), so stars can be exposed about 50% longer on the PC before saturation sets in.

Note that the narrow band filters may be useful when observing very bright targets. For example, stars as bright as V~4.4 can be observed without saturation in F502N using the PC at ATD-GAIN=15 with a 0.11s exposure time.

7.3 Observing Faint Targets Near Bright Objects

The concerns here are similar to those for observing bright targets; saturation and blooming of the bright companion PSF must not impact the faint target. Also, one may need to consider subtracting the PSF of the bright object, and effects which limit the accuracy of that subtraction.

If the bright companion will saturate and bloom, it will be necessary to rotate the CCD so that blooming along the CCD columns does not obliterate the faint target. See Figure 7.9 for an illustration of the bloom directions. It may also be useful to orient the field so that the OTA diffraction spikes from the bright companion (along diagonal lines on the CCDs) avoid the faint target. Table 7.1 summarizes ORIENTs which can be used to avoid CCD blooming tracks and OTA diffraction spikes caused by bright objects. For example, if a faint companion is at PA 60°on the sky relative to a bright companion, it would be advantageous to observe on PC1 with ORIENT= PA + 45° = 105°. Ideally, some range in ORIENT would be specified to ease scheduling, hence "ORIENT=90D TO 120D" might be specified on the Phase II proposal. Note that "ORIENT=270D TO 300D" is also feasible, and should be indicated in the visit level comments.

Table 7.1: ORIENTs for Avoiding Bloom Tracks and Diffraction Spikes. "PA" is the position angle of the faint target relative to the bright object. Note that ORIENT should be between 0D and 360D, so subtract 360°, if necessary. In the proposal these are specified as, e.g., "ORIENT=231D TO 261D".
CCD	ORIENT
PC1	PA+30° to PA+60°, PA+210° to PA+240°
WF2	PA+120° to PA+150°, PA+300° to PA+330°
WF3	PA+30° to PA+60°, PA+210° to PA+240°
WF4	PA+120° to PA+150°, PA+300° to PA+330°

If instead of observing a known companion, one is searching for companions, it is advisable to observe at several ORIENTs so that the CCD bloom track and OTA diffraction spikes will not hide possible companions. For example, three ORIENTs, each separated by 60°, would give good data at all possible companion position angles.

If PSF subtraction will be needed during data analysis, then the PC CCD may have some advantage, since it provides better sampling of fine undulations in the PSF. It may also be useful to obtain observations of a second bright star for PSF calibration, though these may be of limited utility since thermal effects and OTA "breathing" can modify the telescope focus, and hence the PSF, on time scales of less than one hour. Any such PSF star should be similar in color to the target, and should be observed at the same CCD position (within 1") and with the same filter. Sub-pixel dithering may also be useful, so as to improve sampling of the PSF (see Dithering with WFPC2).

Figure 7.2 illustrates the effect of OTA breathing, and periodic focus adjustments, on PSF subtraction. It shows the difference between an "in focus" PSF and one where the OTA secondary mirror has been moved by 5µm. This amount of focus change is comparable to the range of OTA "breathing" effects (time scale <1 hour), and the periodic (semi-annual) focus adjustments of the OTA. Each panel shows a different contrast setting; the percentages indicate the energy per pixel which is plotted as white, expressed as a fraction of the total (un-subtracted) PSF energy. For example, features which are just white in the "0.003%" panel contain 0.003% of the total PSF energy in each pixel. In other words, the feature labeled "a" is, in effect, ~10 magnitudes fainter than the PSF of the bright object, so that it may be very difficult to detect a "real" companion object ~10 magnitudes fainter than the bright object, at this distance from the bright object. In a real PSF subtraction situation, other effects including PSF sampling, noise, and pointing instability would further degrade the subtraction. (The elongated appearance of the residuals in the PSF core is due to astigmatism in PC1).

Table 7.2 gives some quantitative indication of the performance expected for PSF subtractions in the high signal-to-noise limit. It gives the magnitude of "star-like" artifacts remaining in the subtracted image, as a function of distance from the bright object, and magnitude m_bright for the bright object. The right-most column gives an effective magnitude limit imposed by artifacts from the PSF subtraction. These results are derived for the 5µm focus shift described above, and are for PC1 and filter F555W. It may be possible to do somewhat better than these limits by subtracting accurate model PSFs, or by finding an observed PSF with matching focus.

***Table 7.2: Approx. PSF Subtraction Artifact Magnitudes and Magnitude Limits.***
Distance from Bright Object	Effective Magnitude of Subtraction Artifacts	Effective Faint Object Detection Limit (3)
0.1"	m_bright+4.7	m_bright+3.5
0.3"	m_bright+8.6	m_bright+7.4
1"	m_bright+11.4	m_bright+10.2
3"	m_bright+13.2	m_bright+12.0

Results indicate that PSF subtraction and detection of faint objects very close to bright objects can be improved by using a composite PSF from real data, especially dithered data. Table 7.3 indicates limits that may be obtained for well-exposed sources (nominal S/N > 10 for the faint object) where a dithered PSF image has been obtained.

***Table 7.3: Limiting Magnitudes for PSF Subtraction Near Bright Objects.***
Separation in arcsec (on PC)	Limiting m (without PSF subtraction)	Limiting m (with PSF subtraction)
0.15	2.5	5.0
0.25	4.5	6.4
0.4	6.5	7.3
1.0	8.9	10.7
3.0	10.7	12.9

A technique that has been used with some success to search for nearby neighbors of bright stars is to image the source at two different roll angles, and use one observation as the model PSF for the other. In the difference image, the secondary source will appear as a positive residual at one position and a negative residual at a position separated by the change in roll angles. PSF artifacts generally do not depend on roll angle, but rather are fixed with respect to the telescope. Thus, small changes in the PSF between observations will not display the positive or negative signature of a true astrophysical object. Again, it is recommended that the observations at each roll angle be dithered.

Large angle scattering may also impact identification of very faint objects near very bright ones. This scattering appears to occur primarily in the camera relay optics, or in the CCD. Hence, if a faint target is more than ~10" from a bright object (i.e. very highly saturated object), it would be advisable to place the bright object on a different CCD, so as to minimize large angle scattering in the camera containing the faint target. See the section on Large Angle Scattering. Note also that highly saturated PSFs exist for PC1 in filters F439W, F555W, F675W, and F814W, and for F606W on WF3; these may be useful when attempting to subtract the large-angle scattered light. As of this writing TinyTIM does not accurately model the large angle scattering, and should be used with caution when analyzing highly saturated images (Krist 1996). To obtain available PSFs please visit the PSF Library page at:

It is generally unwise to place bright companions or other bright objects just outside the area imaged by the CCDs. The region of the focal plane just outside the CCDs (within about 6" of the CCDs) contains a number of surfaces which can reflect light back onto the CCDs, hence placing bright targets there can have undesired results. Also, the un-imaged "L" shaped region surrounding PC1 should be avoided, since incomplete baffling of the relay optics allows out-of-focus images of objects in this region to fall on the CCDs. Figure 7.3 illustrates various bright object avoidance regions near the WFPC2 field-of-view; the indicated avoidance magnitudes will produce 0.0016 e^- s^-1 pixel^-1 in the stray light pattern for F555W. Figure 7.4 and Figure 7.5 show examples of artifacts which can result from bright stars near the PC1 CCD. The report "A Field Guide to WFPC2 Image Anomalies" (ISR WFPC2 95-06, available on the WFPC2 WWW pages and from (help@stsci.edu) gives more detailed discussions of artifacts associated with bright objects, and their avoidance.

7.4 Cosmic Rays

Cosmic rays will obliterate ~20 pixels per second per CCD. It is imperative that two or more images be obtained at each pointing position, if these artifacts are to be removed from the data. The default action by the Phase II proposal processing software is to split exposures longer than 600s into two nearly equal parts, so as to allow removal of the cosmic ray tracks. The CR-SPLIT and CR-TOLERANCE optional parameters on the Phase II proposal allow observers to adjust this behavior. CR-SPLIT can be set to either DEF (default), NO, or a numeric value (0.0 to 1.0) giving the fraction of the total exposure allotted to the first sub-exposure of the pair. CR-TOLERANCE indicates the spread allowed in dividing the exposure, as a fraction of the total exposure time. For example, the default CR-TOLERANCE=0.2 allows the first sub-exposure to range from 0.3 to 0.7 of the total exposure. Setting CR-TOLERANCE=0 will force equal-length sub-exposures.

The required degree of cosmic-ray avoidance will depend on the science goals of the proposal; observations of a single small target will usually suffer much less impact from cosmic rays than programs needing very "clean" data over a large area. Table 7.4 gives very rough recommendations for the number of sub-exposures for a given total exposure time. Note that splitting into many sub-exposures introduces additional overhead time and will increase the noise for "read noise" limited exposures (usually exposures in UV or narrow band filters), and hence one should not use more sub-exposures than are truly required by the science goals.

***Table 7.4: Recommended Exposure Splittings.***
Total Exposure Time (s)	Rough Recommended Number of Sub-exposures
Total Exposure Time (s)	Programs with Single Small Target	Wide-area Search Programs
<300	1	3
300 - 600	1 or 2	4
600 - 1600	2 or 3	4
1600 - 5000	3	5
5000 - 10000	4	6
>10000	One exposure per orbit (2400s each)

7.5 Choosing Exposure Times

The choice of exposure time generally depends on the signal-to-noise ratio required to meet the science goals. This can be assessed using information in Chapter 6 or plots in Appendix B herein, or by using the on-line WWW Exposure Time Calculator tool.

However, when packing orbits, one must often compromise somewhat and decide which exposures to lengthen or shorten. Table 7.5 may be helpful in this regard. It shows the total time required to execute a single CR-SPLIT=NO exposure, excluding any time needed to change filters.

Note that the most efficient exposure times are those whose length approaches or equals, but does not exceed, an integral number of minutes plus 40s. Figure 7.6 illustrates event timings during a typical 60s WFPC2 exposure, similarly, Figure 7.7 illustrates events during a (more efficient) 100s exposure. (See Overhead Times for more information about exposure timings).

Due to the various overheads, shortening or lengthening an exposure can have unexpected effects on the orbit packing. For example, shortening an exposure from 400s to 350s has no effect on orbit packing; they both require 9 minutes to execute (CLOCKS=NO, the default setting). On the other hand, shortening an exposure from 180s to 160s trims the execution time by 2 minutes (again CLOCKS=NO, the default setting).

CLOCKS=YES may have some advantage in a long series of exposures whose lengths are 180s or somewhat greater. Each savings of 1 minute can add up to a few more exposures per orbit. The down side is that most calibrations are derived for exposures with CLOCKS=NO, so the calibration may be slightly compromised. The largest calibration error is expected to occur in the dark current, where there may be a slight increase near the top and bottom of each CCD. In many situations this error may be acceptable, such as a small target near a CCD center, or broad band filter images where the sky completely dominates the dark current. CLOCKS=YES will have more impact on calibration of narrow filters, or situations requiring an extremely flat background. (Also, see Serial Clocks for discussion of exposure time anomalies associated with CLOCKS=YES, though these are most important for exposures <30s.)

An exposure with CR-SPLIT=YES would simply require the total time for each sub-exposure as given by Table 7.5, again, plus any time needed to change filter before the first exposure. However, the default CR-SPLITting allows schedulers some latitude in dividing the exposures (CR-TOLERANCE=0.2 is the default) so the exact overheads are unpredictable. For example, a 700s exposure with CR-SPLIT=0.5 (the default) could be split into a pair of 350s exposures totaling 18 minutes, or a 300s and 400s exposure totaling 17 minutes.

Table 7.5: Basic Time to Execute Single Non-CR-SPLIT Exposure. This includes time to prep the CCD, execute the exposure, and read out the CCDs. Times needed to change filter (1 minute), or insert a second filter (1 minute), are excluded. See Overhead Times for more discussion and other overheads.
Exposure Time (s)	Total Execution Time (min.)
Exposure Time (s)	CLOCKS=NO (default)	CLOCKS=YES
0.11 to 30	2	(not recommended)
35, 40	2	2
50,60,70,80,100	3	3
120,140,160	4	4
180,200	6	5
230,260	7	6
300	8	7
350,400	9	8
500	11	10
600	13	12
700	14	13
800	16	15
900	18	17
1000	19	18
1100	21	20
1200	23	22

7.6 Dithering with WFPC2

Dithering is the technique of displacing the telescope between observations either on integral pixel scales (to assist in removing chip blemishes such as hot pixels) or on sub-pixel scales (to improve sampling and thus produce a higher-quality final image). Here we briefly discuss observation and data analysis for dithered data.

7.6.1 Dither Strategies

There is no single observing strategy that is entirely satisfactory in all circumstances for WFPC2. One must consider cosmic rays, hot pixels (i.e. pixels with high, time variable dark count), spatial undersampling of the image, and large-scale irregularities such as the few arcsecond wide region where the CCDs adjoin. One strategy that can be used to minimize the effects of undersampling and to reduce the effects of hot pixels and imperfect flat fields is to dither, that is, to offset the telescope by either integer-pixel or sub-pixel steps. The best choice for the number and size of the dithers depends on the amount of time available and the goals of the project. In the following we will address a few issues related to dithering:

The recommended way to schedule dithers is to specify dither patterns WFPC2-LINE (e.g. for two-point diagonal dithers) or WFPC2-BOX (for four-point dithers). An alternative approach is to use POS TARGs. Note that when the WF3 is specified as an aperture, the POS TARG axes run exactly along the WF3 rows and columns. For the other chips, they only run approximately along the rows and columns due to the small amount of rotation between CCDs. For small dithers (less than a few pixels) these rotations are unimportant.

Some specific offsets allow one to shift by convenient amounts both the PC and the WFC chips. For instance an offset of 0.5" is equivalent to 5 WFC pixels and 11 PC pixels. Likewise, the default WFPC2-LINE spacing of 0.3535" along the diagonal is equivalent to shifts of (2.5,2.5) pixels for the WFC and (5.5,5.5) pixels for the PC.

Dithers larger than a few pixels will incur errors due to the camera geometric distortion which increases toward the CCD corners and alters the image scale by about 2% at the corners. Hence a 1.993" offset will be 20.3 WF pixels at the field center, but suffer a 0.4 pixel error at the CCD corners. Large dithers may also occasionally require a different set of guide stars for each pointing, thus greatly reducing the expected pointing accuracy (accuracy only ~1" due to guide star catalogue).

The most up-to-date information about dither strategies and related issues can be found on the general WFPC2 dither web page:

7.6.2 Analysis of Dithered Data

The software we recommend for combining dithered data is known as "MultiDrizzle" (Koekemoer, et al. 2002), which is based on the "drizzle" program (Fruchter and Hook 2002). This method has been incorporated into the IRAF/STSDAS dither package, and allows effective cosmic ray removal from dithered data.

In order to help users reduce dithered images, we have prepared the HST Dither Handbook (Koekemoer et al. 2002), available from the above WFPC2 dither web site. This document gives a general outline of the reduction of dithered images and provides step-by-step instructions for six real-life examples that cover a range of characteristics users might encounter in their observations. The data and scripts needed to reproduce the examples are also available via the same URL. (This handbook expands upon the original Drizzling Cookbook by Gonzaga et al. 1998.)

Despite all the improvements in the combination of dithered images, users should be mindful of the following cautionary notes:

7.7 Pointing Accuracy

Some WFPC2 programs have critical target positioning constraints (i.e. the target must be as close as possible to a specified aperture). A sure way to meet such requirements is to include an interactive acquisition. However, INT ACQs are costly in terms of allotted orbits. A variation of the Reuse Target Offset (RTO) capability can be used to acquire and position a target in the WFPC2 FOV. However, the user must request an additional orbit for the acquisition. The first orbit is used for the acquisition and the second orbit for the science observations.

7.7.1 Absolute Pointing Accuracy

We have looked carefully at a sequence of images to assess the absolute pointing performance that HST delivers to WFPC2. The apertures used in the observations studied were either PC1, PC1-FIX, or WF2. The observed positions of stars on WFPC2 images were measured and compared with the proposed coordinates and apertures. Where necessary, coordinate and proper motion errors were accounted for (with the assumption that SAO catalog coordinates are exact - they form the astrometric basis for the guide star coordinate system). The typical residual pointing error is 0.86", with 1.84" being the largest error seen. This study did bring out several easy-to-make target coordinate errors (which we corrected in the analysis, but which frequently dominated the pointing error), so we discuss these first.

In a number of cases studied, the proposal coordinates were from the printed version of the Yale Bright Star Catalog. One problem is that the equinox 2000 positions in the BSC are given in the FK4 (Besselian) reference system. The proposal system assumes that equinox 2000 and later coordinates are in the FK5 (Julian) reference frame, and that earlier ones are in the FK4 frame. This can be overridden by specifying B2000 instead of J2000 for the equinox in the proposal. The latest digital version of the BSC (BSC5) is in J2000. The 1950 edition of the SAO catalog is in B1950 (FK4), and a digital version is available for J2000 (FK5). An error of up to 1.5" can result from assuming BSC positions are J2000 instead of B2000 in the proposal.

Another common problem with target coordinates is that they lack precision. For example, in the BSC, RA is given to the nearest 0.1^s and DEC to the nearest arcsecond. This can cause an error of up to 0.75" in RA and 0.5" in DEC. The SAO coordinates have higher precision, 0.001^s in RA and 0.01" in DEC, and should be used when possible.

A common error source is not specifying proper motion or specifying it in the wrong units. It is critical to follow the latest version of the proposal instructions on this. Even small proper motions are significant at the resolution of HST images.

Residual pointing errors (after coordinate errors and aperture location changes) range from 0.26" to 1.84". The average is 0.93" and the median is 0.86". There are no obvious trends in any coordinate system. These are errors which cannot be accounted for by a proposer, being due to guide star position errors, FGS alignment uncertainties, and residual aperture location errors. Using Guide Star Catalog positions may help reduce the error between target and guide stars. Most of the targets used in this study were too bright to have true Guide Star Catalog positions.

In summary, a target with good coordinates (and proper motion) referenced to the SAO catalog can typically be placed within 0.9" of a specified aperture. However, errors of around 1.5" occasionally happen.

7.7.2 Updates to Aperture / Coordinate Systems

On 11 April, 1994, an update was made to the spacecraft database which tells HST where to place targets relative to the FGSs. This update affected both the location of targets in the WFPC2 field-of-view, and the position reference frame in the image headers. The nominal (or intended) pixel locations of the apertures in the WFPC2 focal plane did not change. Only the (V2,V3) coordinates of the apertures changed, as their locations relative to the FGSs became better known. For example, PC1 and PC1-FIX are designated to be at pixel (420,424.5). Before April 1994, this aperture was thought to be at (V2,V3) =(4.95",-30.77"), which, using the most current information, was actually located at pixel (X,Y) = (459.8,377.3). Since April 1994, the aperture in the spacecraft's database has been at (V2,V3) = (1.87",-30.96") or, assuming the current best estimate is exactly correct, at (X,Y) = (414.9,428.1). Thus, for the same coordinates and aperture, the pixel position of a target in an image taken before April 1994, could be nearly 3" different from its position in later images, due to aperture updates. Similar corrections apply to all WFPC2 data taken before this date.

This update also affects the position information placed in the image headers, which maps sky coordinates onto each individual CCD. Observations taken before April 11, 1994, have preliminary plate scales, rotations, and reference pixel locations in their image headers. Thus, the sky coordinates associated with a given pixel will be different for otherwise identical images taken before and after April 11, 1994, due to improvements in the aperture locations. The change is primarily an approximate 3" shift, as well as a small rotation. There is a 0.8° rotation for WF2, and smaller rotations for the other chips (0.28° in PC1, 0.46° in WF3, and 0.06° in WF4). We note that the On-The-Fly Calibration System initiated in 2000 does not correct for these offsets, since the pointing information is set upstream of the pipeline calibration; the On-The-Fly Reprocessing System installed in May 2001 does, however, correct the pointing offsets.

The STSDAS tasks METRIC/INVMETRIC and WMOSAIC use this header information; hence, images taken before April 11, 1994, required header updates in order for these tasks to produce optimum results. In this situation, observers were advised to run the STSDAS task UCHCOORD, to update the headers, prior to running METRIC/INVMETRIC and/or WMOSAIC.

The On-The-Fly Calibration System (OTFC), in place from Dec. 1999 to May 2001, did not correct for these offsets. Observers submitting requests to the archive prior to May 16, 2001 received data processed through OTFC; this data would benefit from running UCHCOORD.

As of May 16, 2001, however, the On-The-Fly Reprocessing System (OTFR) is in place and OTFR data does contain the most up-to-date header information.

	*UCHCOORD should not be run on OTFR data.*

OTFR data can be identified by the presence of the keyword PROCTIME in the header. Please see On-The-Fly Reprocessing Systems for more details on OTFC, OTFR, and the use of UCHCOORD.

We also note, that in April and May 1996, two updates were made to the (V2,V3) coordinate system. This update should not affect observers. The purpose was to remove a slow drift in the position of WFPC2 in the HST focal plane; the largest change was 0.6". (See Table 3.15 on page 75 for details.) An additional update of 0.2" was made on December 1, 1997. All the apertures are now thought to be correct to within 0.3", and future updates should be small. Please also see the section on the ORIENT Anomaly.

7.7.3 Pointing Repeatability

The Hubble Deep Field (HDF) afforded an opportunity to study the repeatability of pointing over many images and acquisitions of the same field. The pointing appears to have been stable to better than 5 mas accuracy while taking many images of the same field without interruption over several orbits. The accuracy for full-up acquisition of the same field after slewing to other targets appeared to be ~10 mas typically, with occasional 20 mas errors seen. However, a few large errors were seen; in about 1 in 100 acquisitions the FGSs locked-up incorrectly resulting in a ~1" error.

Other programs report similar 3 mas pointing accuracy if simple re-acquisitions are done between orbits. Approximately once per day a "full-up" acquisition is usually required (for engineering reasons) where the dominant FGS is fixed in position, but the sub-dominant FGS performs a spiral search for the guidestar and tracks wherever the star is found. On rare occasion these full-up acquisitions produce position errors of several hundred mas, and field rotations of up to ~0.1°, relative to previous images of the same field. This may impact long sequences of exposures requiring half a day or more to execute.

7.7.4 Tracking Modes

Two guiding modes are available: Gyro Hold, and Fine Lock. Fine Lock (PCS MODE FINE) is used by default, since use of Coarse Track may be harmful to the Fine Guidance Sensors. Use of Gyro Hold (PCS MODE GYRO) is not generally recommended, even for snapshot (SNAP) observations, since the pointing accuracy is only 14". Also the drift rate is 0.0014" s^-1 so exposures >100s can result in smeared images. However, if the reduced pointing accuracy can be tolerated, and the exposures are only a few seconds or less, Gyro Hold can give a significant savings in the target acquisition overhead time.

7.8 CCD Position and Orientation on Sky

During observation the target is placed at the aperture (PC1, WF2, WFALL, etc.) specified on the Phase II proposal. Locations of the principal apertures are shown in Figure 7.9 (Table 3.14 on page 74 gives a complete list of apertures; the (V2,V3) system here is post 1996 day 127).

The POS TARG special requirement can be used when a position offset is needed. The target is positioned with offset "POS TARG x,y", measured in arcseconds, from the specified aperture. The approximate directions (within 1°) of the POS TARG offsets are shown in Figure 7.9. The exact directions of the offsets are parallel to the rows and columns of the CCD on which the aperture is specified. There are small rotations (few tenths of a degree) between the CCDs. (For detailed information see "Dithering: Relationship Between POS TARGs and CCD Rows/Columns" obtainable from the WFPC2 WWW pages or help@stsci.edu.)

It is often useful to explicitly specify the desired rotation of the WFPC2 field-of-view on the sky. This is specified in the Phase II proposal using the ORIENT special requirement. It is defined as the PA (measured from North through East) of the +U3 axis on the sky. Figure 7.9 shows the CCD orientation and aperture locations relative to the U3 axis.

In effect, the sequence of events is to first move the target to the desired aperture, then offset by any specified POS TARG from the aperture, and finally to rotate the target "in place" on the CCDs to the desired ORIENT.

Observers should try to specify all possible ORIENTs which would give the desired data, since having a range of values, or several ranges, will make the observation much easier to schedule. Often two ORIENTs separated by 180° will both give useful data. Sometimes ORIENTs separated by 90° will also give similar data.

Another way to select the ORIENT, is to place Figure 7.9 on an image of the target, shift and rotate to get the desired alignment, and then simply measure the position angle of the +U3 axis relative to North.

Note that the +V3 axis is quite different from the +U3 axis. They are exactly parallel, but oppositely directed. The +U3 axis is used for specifying orientation (ORIENT) in the proposal, while the +V3 axis is used in the data headers to indicate field orientation. Data header keyword PA_V3 gives the position angle of the +V3 axis on the sky.

We now give two examples of how the POS TARG and ORIENT special requirements might be used. The first example (Figure 7.10) shows placement of a 100" long jet along the CCD diagonals in PC1 and WF3 (i.e. along the -U3 direction). The coordinates of the nucleus are given on the proposal. Aperture PC1 together with POS TARG +10, +10 are used to place the nucleus near the outer corner of PC1. We want to rotate the WFPC2 field-of-view about the nucleus so the jet is diagonal on PC1 and WF3. We can thus compute the desired orientation as


ORIENT = (source PA on sky) + (desired source angle in field-of-view measured CW from +U3 axis) = 290° + 180° = 470° - 360° = 110°

On the Phase II proposal we would allow some range in the angle (to ease scheduling), hence "ORIENT 105D TO 115D" might be specified.

The second example (Figure 7.11) shows placement of a galaxy across WF2 and WF3, with the nucleus on WF3 safely away from the vignetted region. Aperture WF3 together with POS TARG +20, 0 is used to place the nucleus near the outer edge of WF3. We want to rotate the WFPC2 field-of-view about the nucleus so the galaxy's major axis is across WF2 and WF3. We can thus compute the desired orientation as


ORIENT = (source PA on sky) + (desired source angle in field-of-view measured CW from +U3 axis) = 60° + 315° = 375° - 360° = 15°

On the Phase II proposal we would again allow some range in the angle (to ease scheduling), hence "ORIENT 5D TO 25D" would be specified. Note that "ORIENT 185D TO 205D" is also feasible, and should be indicated in the visit level comments. Note also, that WF3 and WF4 could be used with either "ORIENT 95D TO 115D" or "ORIENT 275D TO 295D".

7.8.1 Software to Aid ORIENT Selection

The Visual Target Tuner (VTT) allows observers to select ORIENTs via a graphical user interface. The VTT will display Digitized Sky Survey images, NASA/IPAC Extragalactic Database (NED) images, HST images, or other suitable FITS images. It will then superpose HST instrument apertures and allow observers to manipulate their ORIENT and position. For more information see http://apt.stsci.edu/vtt/.

7.8.2 ORIENT Anomaly

We note that a minor anomaly was discovered in the data header values pertaining to image orientation (i.e. rotation about the target aperture) for data taken prior to September 15, 1997. Specifically the header keywords PA_V3 and ORIENTAT were affected. During long visits their values were incremented by up to 0.05 degree per hour whenever the telescope pointing was changed, when in fact these header values should have remained fixed. Observers requiring highly accurate image orientations should check values in the so-called jitter files (*jit.fits and *jif.fits), which were not affected by the bug. Data extracted from the HST archive using the On-The-Fly Reprocessing system implemented in mid-2001 is automatically corrected for this problem.

7.9 Polarization Observations

Polarization observations require three or more images with the polarizing filter spanning a large range of position angles on the sky. For WFPC2, this may be achieved by using different quads of the polarizing filter (each quad being oriented 45° to the others), by rotating the spacecraft though different angles, or by a combination of these methods. Rotating the spacecraft through use of ORIENTs provides the simplest calibration, as a single polaroid can be used for all images. However, in practice, it will be the most difficult method to schedule. For further information see Biretta and Sparks (1995, ISR WFPC2 95-01).

We note that WFPC2 has significant instrumental polarizations which will make measurements on targets with less than 3% polarization difficult. The pick-off mirror introduces about 6% instrumental polarization. Furthermore, the pair of mirrors in the calibration channel, which is used to generate the polarizer flat fields, introduces ~12% polarization. In principle these effects can be calibrated out, but this has yet to be demonstrated.

The polarizers are most effective in the range from 3000Å to 6500Å; this corresponds roughly to filters in the range F255W to F675W. At shorter wavelengths the transmission decreases sharply, and at longer wavelengths they cease to polarize the incoming light.

7.10 Observing with Linear Ramp Filters

The Linear Ramp Filters (LRFs) provide a narrow band (

) imaging capability which is continuously tunable from 3710Å to 9762Å. These are essentially a collection of narrow band interference filters whose central wavelength varies with position on the filter glass. The filter and aperture should be specified as LRF on the Phase II proposal, and the desired central wavelength should also be specified. The HST scheduling software will then select the target position so as to provide the desired wavelength.

It is possible to use POS-TARGs with LRF observations; the offsets are made from the default pointing for the specified wavelength. Observers should be mindful that the unvignetted field-of-view has a minimum size of ~10" in diameter, so that only small POS-TARGs (<4") should be used.

While it is recommended that observers assume a 10" diameter field-of-view when using the LRFs, larger elongated (e.g. 15" x 10") targets can sometimes be accommodated by placing the target's major axis along the direction of the wavelength variation on the filter. This will result in a small reduction in throughput (i.e. small central wavelength offset) at the outer edges of the target. However, placing targets outside the central 10" of each ramp is strongly discouraged; outside the central 10" width the light will pass through more than one ramp segment, hence mixing light from different wavelengths, and making the data very difficult to calibrate. (See "Linear Ramp Filters" on page 48 for further details on LRFs.)

A common situation is one in which observers desire to make observations through an LRF filter, and then repeat the observation in a standard broad or narrow band filter at the same position on the CCD. The LRF Calculator Tool, available on the WFPC2 WWW pages, will tell observers the aperture (PC1-FIX, WF2-FIX, etc.) and POS-TARG for any wavelength setting of the LRFs. Observers merely need to use this same aperture and POS-TARG for the exposure through the other filter. If it is necessary to calculate the POS-TARG manually, one can do this using the information in Table 3.7 on page 58, Table 3.8 on page 60, Table 3.14 on page 74, and Figure 7.9. For example an LRF observation at 5034Å would be made on WF2 at pixel (673.4, 235.7) (from interpolation by wavelength between X1 and X2, and between Y1 and Y2 in Table 3.7 on page 58). These offsets are referred to the WF2-FIX aperture which is located (Table 3.14 on page 74) at pixel (423.5,414). From Figure 7.9 we can deduce that pixel X direction is parallel to POS-TARG "+Y" on WF2, and that pixel Y direction runs in the POS-TARG "-X" direction. Using the pixel scale in Optical Distortion, we have

We note that analysis of FR533N VISFLAT images has revealed an apparently randomly occurring offset in the filter position (Gonzaga et al. 2001, WFPC2 ISR 01-04). This anomalous offset corresponds to one step in the filter wheel rotation, or about 0.5 degrees. We expect no significant impact on point-source observations; any photometric effect is less than 1%. But caution needs to be exercised for extended sources greater than about 5 arcseconds. (A cursory check of several other filters on other filter wheels shows no similar problems.) Figure 7.12 shows throughput plotted against CCD pixels in the direction of the anomalous offset/rotation. The two curves in each plot show the throughput effect of the filter offset. Several points in the wavelength mapping (from actual GRW+70D5824 observations in proposals 6939, 8054, and 8454) are indicated for illustrative purposes. At this time, the source of this anomaly, whether it is mechanical or due to a software error, is not known. This anomaly was investigated further as part of the WFPC2 Cycle 10 Calibration Plan.

7.11 Emission Line Observations of Galaxy Nuclei

Saturation is a common problem for narrow band filter images of galaxy nuclei. Often the surface brightness of the emission line is estimated from ground based images with 1" resolution; sometimes line fluxes are quoted for apertures several arcseconds in radius. However, at HST resolution, much of this flux may occur in a single unresolved spot at the galaxy nucleus, thus leading to saturated images. Observers should be cautious and consider this possibility when estimating exposure times.

7.12 Two-Gyro Mode

At some future date HST may be operated with only two gyros, hence causing additional spacecraft jitter and degradation of the effective PSF. Other potential impacts of operating in two-gyro mode are decreased visibility periods between occultations, and restricted target availability during the year. It is also possible that the pointing stability between exposures may be slightly poorer.


	*Please see the `Two-Gyro Mode Handbook` for additional discussion.*

Wide Field and Planetary Camera 2 Instrument Handbook for Cycle 14

Chapter 7:
Observation Strategies

7.1 Observing Faint Targets

7.2 Observing Bright Targets

7.3 Observing Faint Targets Near Bright Objects

7.4 Cosmic Rays

7.5 Choosing Exposure Times

7.6 Dithering with WFPC2

7.6.1 Dither Strategies

7.6.2 Analysis of Dithered Data

7.7 Pointing Accuracy

7.7.1 Absolute Pointing Accuracy

7.7.2 Updates to Aperture / Coordinate Systems

7.7.3 Pointing Repeatability

7.7.4 Tracking Modes

7.8 CCD Position and Orientation on Sky

7.8.1 Software to Aid ORIENT Selection

7.8.2 ORIENT Anomaly

7.9 Polarization Observations

7.10 Observing with Linear Ramp Filters

7.11 Emission Line Observations of Galaxy Nuclei

7.12 Two-Gyro Mode

Wide Field and Planetary Camera 2 Instrument Handbook for Cycle 14

Chapter 7: Observation Strategies

7.1 Observing Faint Targets

7.2 Observing Bright Targets

7.3 Observing Faint Targets Near Bright Objects

7.4 Cosmic Rays

7.5 Choosing Exposure Times

7.6 Dithering with WFPC2

7.6.1 Dither Strategies

7.6.2 Analysis of Dithered Data

7.7 Pointing Accuracy

7.7.1 Absolute Pointing Accuracy

7.7.2 Updates to Aperture / Coordinate Systems

7.7.3 Pointing Repeatability

7.7.4 Tracking Modes

7.8 CCD Position and Orientation on Sky

7.8.1 Software to Aid ORIENT Selection

7.8.2 ORIENT Anomaly

7.9 Polarization Observations

7.10 Observing with Linear Ramp Filters

7.11 Emission Line Observations of Galaxy Nuclei

7.12 Two-Gyro Mode

Chapter 7:
Observation Strategies