7.10	IR Observing Strategies

7.10.1 Dithering Strategies

7.10.2 Parallel Observations

7.10.3 Exposure Strategies

7.10.4 Spatial Scans

7.10.5 PSF Subtraction

7.10.1 Dithering Strategies

For imaging programs, STScI generally recommends that observers employ dithering patterns. Dithering refers to the procedure of moving the telescope by small angle offsets between individual exposures on a target. The resulting images are subsequently combined in the pipeline or by the observer using software such as AstroDrizzle. (See the DrizzlePac documentation.)

Dithering is used to improve image quality and resolution. By combining multiple images of a target at slightly different positions on the detector, one can compensate for detector artifacts (blemishes, dead pixels, hot pixels, transient bad pixels, and plate-scale irregularities) that may not be completely corrected by application of the calibration reference files. Combining images, whether dithered or not, can also remove any residual cosmic ray flux that has not been well removed by the up-the-ramp fitting procedure used to produce flt images (see Section 7.7.2 and Appendix E:Reduction and Calibration of WFC3 Data). Effective resolution can be improved by combining images made with sub-pixel offsets designed to better sample the PSF. This is especially important for WFC3/IR, because the PSF is undersampled by about a factor of 2 (see Table 7.5).

Larger offsets are used to mosaic a region of sky larger than the detector field of view. (Large offsets can also be used for “chopping” to sample the thermal background. This has been recommended for NICMOS exposures at wavelengths longer than 1.7 microns, where the telescope thermal background becomes increasingly dominant, but the thermal background is not a problem for WFC3/IR). In WFC3, all offsets must be accomplished by moving the telescope (whereas in NICMOS it was also possible to move the Field Offset Mirror). Dithers must be contained within a diameter ~130 arcsec or less (depending on the availability of guide stars in the region) to use the same guide stars for all exposures. The rms pointing repeatability is significantly less accurate if different guide stars are used for some exposures. (See Appendix B of the DrizzlePac Handbook.). Mosaic steps and small dither steps are often combined to increase the sky coverage while also increasing resolution and removing artifacts. (See Section 6.11.1 for a discussion of the effect of geometric distortion on PSF sampling for mosaic steps).

The set of Pattern Parameters in the observing proposal provides a convenient means for specifying the desired pattern of offsets. The pre-defined mosaic and dither patterns that have been implemented in APT to meet many of the needs outlined above are described in detail in the Phase II Proposal Instructions. The WFC3 patterns in effect in APT at the time of publication of this Handbook are summarized in Appendix C:Dithering and Mosaicking. Observers can define their own patters to tailor them to the amount of allocated observing time and the desired science goals of the program. Alternatively, they can use POS TARGs to implement dither steps (Section 7.4.3). Observers should note that thermally driven drift of the image on the detector, occasionally larger than 0.15 pixels in two orbits, will limit the accuracy of execution of dither patterns. (WFC3 ISR 2009-32) Additional information on dither strategies can be found in WFC3 ISR 2010-09.

7.10.2 Parallel Observations

Parallel observations, i.e., the simultaneous use of WFC3 with one or more other HST instruments, are the same for the IR channel as for the UVIS channel, previously described in Section 6.11.2.

7.10.3 Exposure Strategies

Given the variety of requirements of the scientific programs that are being executed with WFC3/IR, it is impossible to establish a single optimum observing strategy. In this section we therefore provide a few examples after guiding the reader through the main constraints that should be taken into account when designing an observation:

•	Integrate long enough to be limited by background emission and not read noise. Dark current is rarely the limiting factor.

•	Dither enough so that resolution can be restored to the diffraction limit and bad pixels and cosmic-ray impacts can be removed, while maintaining a homogeneous S/N ratio across the field.

•	Split the MULTIACCUM ramps into as many reads as possible for readout noise suppression.

These constraints put contradictory requirements on the ideal observing strategy. It is clear that, given a certain amount of total observing time, the requirement of long integrations for background limited performance is incompatible with a large number of dithering positions. Also, to split ramps for readout noise suppression decreases the observing efficiency, with a negative impact on the signal to noise ratio. Because the background seen by each pixel depends on the filter (Section 7.9.5), the optimal compromise must be determined on a case-by-case basis.

In this regard, it is useful to consider Table 7.11, which summarizes the total background seen by a pixel, including sky, telescope, and nominal dark current, and the time needed to reach 400 e–/pixel of accumulated signal, corresponding to 20 e–/pixel of Poisson-distributed background noise. This last value, higher than the expected readout noise of ~12 electrons after 16 reads, is used here to set the threshold for background-limited performance. The passage from readout-limited performance to background-limited performance can be regarded as the optimal exposure time for that given filter, in the sense that it allows for the largest number of dithered images without significant loss of S/N ratio (for a given total exposure time, i.e., neglecting overheads). For faint sources, the optimal integration time strongly depends on the background (zodiacal, Earth-shine thermal, and dark current) in each filter, ranging from just 220 s for the F110W filter to 2700 s for some of the narrow-band filters.

The optimal integration time needed to reach background-limited performance (see Table 7.11) can be compared with the integration times of the sampling sequences from Table 7.8. Table 7.12 synthesizes the results, showing for each filter which ramp (SPARS, STEP) most closely matches the optimal integration times for NSAMP=15.

Table 7.11: Background (e–/pix/s) levels at the WFC3/IR detector. The columns show, from left to right: a) filter name; b) thermal background from the telescope and instrument; c) zodiacal background; d) earth-shine background; e) dark current; f) total background; g) integration time needed to reach background-limited performance, set at an equivalent readout noise of 20 electrons.

Filter	Thermal	Zodiacal	Earth- shine	Dark Current	Total	Optimal Integration Time (sec)
F105W	0.051	0.774	0.238	0.048	1.111	360
F110W	0.052	1.313	0.391	0.048	1.804	222
F125W	0.052	0.786	0.226	0.048	1.112	360
F140W	0.070	0.968	0.267	0.048	1.353	296
F160W	0.134	0.601	0.159	0.048	0.942	425
F098M	0.051	0.444	0.140	0.048	0.683	586
F127M	0.051	0.183	0.052	0.048	0.334	1198
F139M	0.052	0.159	0.044	0.048	0.303	1320
F153M	0.060	0.153	0.041	0.048	0.302	1325
F126N	0.051	0.037	0.011	0.048	0.147	2721
F128N	0.051	0.040	0.011	0.048	0.150	2667
F130N	0.051	0.041	0.011	0.048	0.151	2649
F132N	0.051	0.039	0.011	0.048	0.149	2685
F164N	0.065	0.036	0.009	0.048	0.158	2532
F167N	0.071	0.035	0.009	0.048	0.163	2454

The selection of which sample sequence type (RAPID, SPARS, STEP; Section 7.7.3) must take into account the science goals and the restrictions placed on their use. Here are some factors to consider when selecting a sample sequence:

•	The RAPID ramp is a uniform sequence of short exposures. With its relatively short maximum exposure time, is suitable for a target consisting of bright objects that would saturate after a few reads in the other sequences. It is not appropriate for background-limited performance.

•	SPARS ramps, with their uniform sampling, provide the most robust rejection of cosmic-ray events, and can be trimmed by removing a few of the final reads to fine-tune the integration time with little degradation of the achieved readout noise. Thus they are considered the standard sampling mode.

•

STEP ramps are preferable where large dynamic range is needed; e.g., for photometry of stellar clusters. These ramps begin with a sequence of four uniform (RAPID) reads and end with a sequence of much longer uniform reads. The transition between the two uniform read rates is provided by a short sequence of logarithmically increasing read times. This design provides for correction of any non-linearities early in the exposure and allows for increased dynamic range for both bright and faint targets.

Finally, the selection of a given sample sequence type should also be made in conjunction with the number of samples (nsamp) that will be used to achieve the desired total exposure time for the observation. Long exposures should in general use a minimum of 5-6 samples in order to allow for reliable CR rejection and to allow for at least a few unsaturated samples of bright targets in the field. For very faint targets in read-noise limited exposures, a larger number of samples will result in greater reduction of the net read noise and a more reliable fit to sources with low signal. Short exposures of bright targets, on the other hand, can get by with fewer samples. This is especially true, for example, for the direct images that accompany grism observations. Because the purpose of the direct image is to simply measure the location of sources - as opposed to accurate photometry - they can reliably use an nsamp of only 2 or 3.

Table 7.12: Optimal exposure time needed to reach background-limited performance (see Table 7.11) for each WFC3/IR filter, along with the NSAMP=15 sequences that provide the closest match. The benefits and disadvantages of each sequence type are discussed in the accompanying text.

Filter	Optimal exposure time (sec)	SPARS	STEP
F105W	360	SPARS25	STEP50
F110W	222	SPARS25	STEP25
F125W	360	SPARS25	STEP50
F140W	296	SPARS25	STEP25
F160W	425	SPARS25	STEP50
F098M	586	SPARS50	STEP50
F127M	1198	SPARS100	STEP200
F139M	1320	SPARS100	STEP200
F153M	1325	SPARS100	STEP200
F126N	2721	SPARS200	STEP400
F128N	2667	SPARS200	STEP400
F130N	2649	SPARS200	STEP400
F132N	2685	SPARS200	STEP400
F164N	2532	SPARS200	STEP400
F167N	2454	SPARS200	STEP400

7.10.4 Spatial Scans

Spatial scanning is available with either WFC3 detector, UVIS or IR. Conceptually producing star trails on the IR detector is the same as producing star trails on the UVIS detector, with a few differences discussed in WFC3 ISR 2012-08. This document is recommended to anyone preparing a phase II proposal that uses spatial scans for any purpose. Spatial scans also are discussed elsewhere in this Handbook, specifically for UVIS imaging (Section 6.11.3) and IR slitless spectroscopy (Section 8.6). The former section describes star trails and the latter section describes spectra trailed perpendicular to the dispersion direction.

7.10.5 PSF Subtraction

IR imaging has been shown to be highly effective in detecting faint point sources near bright point sources (WFC3 ISR 2011-07). In this study, deep dithered exposures of a star were made at a variety of roll angles. Unsaturated exposures of a star, scaled down in flux to simulate faint companions of various magnitudes, were added to the deep exposures. The faintness of the companion that can be detected at a certain separation from the bright star depends on the degree of sophistication used to generate a reference image of the PSF to subtract from each set of dithered exposures. For a separation of 1.0 arcsec, five sigma detections could be made fairly easily for companions 8 or 9 magnitudes fainter than the bright star, and companions more than 12 magnitudes fainter than the bright star could be detected at separations of a few arcsec. Substantial improvements in detectability at separations less than about 2 arcsec could be made using the methodology described in the ISR to generate the reference PSF.

If observers want to use stellar images to subtract the PSF from a target comprised of a point source and an extended source to detect or measure the extended source, they should keep several points in mind:

•	IR pixels undersample the PSF (Section 7.6), so the stellar and target exposures should be dithered to produce good sampling of the PSF.

•	Position drift and reacquisition errors can broaden the PSF (WFC3 ISR 2009-32, WFC3 ISR 2012-14).

•

If a single guide star is used for a visit, the roll angle drift results in position drift of the PSF on the detector, typically ~15 mas per orbit (~ 0.1 IR pix in ~45 minutes) at the center of the IR detector, but on rare occasions as large as ~50 mas per orbit (Section B.2 in the DrizzlePac Handbook).

•	The characteristics of the PSF depend on the focus, which generally changes measurably during an orbit; its range in a particular orbit will not be known in advance (WFC3 ISR 2012-14).

•	The characteristics of the PSF vary with location on the detector (e.g., see ACS ISR 2003-06).

•	If exposures long enough to have good signal-to-noise to the desired radius will saturate the central pixels, one should consider making shorter exposures to avoid the effects of persistence. (Section 7.9.4, Appendix D:Bright-Object Constraints and Image Persistence)

While Tiny Tim modeling is available for the WFC3 IR detector, it has not been optimized to reproduce observed PSFs. Progress has been made in understanding the short-comings in the model implemented in version 7.4 (WFC3 ISR 2012-13).