Larger offsets, up to sizes approaching the detector’s field of view, can also be used to create mosaics. However, as a result of geometric distortion (Appendix B:
), some objects shift by an integer number of rows (or columns), while others shift by an integer plus some fraction of a pixel. The PSF is not resampled in that dimension in the former case, but is resampled in the latter case. Where the exposures overlap, the PSF is thus better sampled for some objects than for others. If PSF sampling is important, a combination of mosaic steps and small dither steps should therefore be used. Note that, in practice, mosaic steps 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
Dither strategies for WFC3 are further discussed in WFC3 ISR 2010-09
, which provides a decision tree for selecting patterns and combining them with subpatterns.
While the design of WFC3 precludes the simultaneous use of both the UVIS and IR channel, it is possible to use one or more of the other HST
instruments in parallel with WFC3. Since each instrument covers a different location in the HST
focal plane (see Figure 2.2
), parallel observations typically sample an area of sky several arc minutes away from the WFC3 target. For extended targets such as nearby galaxies, parallel observations may be able to sample adjacent regions of the primary target. In other cases, the parallel observations may look at essentially random areas of sky.
For processing and scheduling purposes, HST
parallel observations are divided into two groups: coordinated and pure.
A coordinated parallel
is an observation directly related to (i.e., coordinated with) a specific primary observation, such as in the extended galaxy example above. A pure parallel
is an observation typically unrelated to the primary observation, for example, parallel imaging scheduled during long spectroscopic observations. The primary restriction on parallel observations, both coordinated and pure, is that they must not interfere with the primary observations: they may not cause the primary observations to be shortened; and they must not cause the stored-command capacity and data-volume limits to be exceeded. The proposal software (APT) enforces these rules and notifies the observer when a specified parallel is not permitted.
In order to prolong the life of the HST
transmitters, the number of parallels acquired during each proposal cycle is limited. Proposers must provide clear and strong justification in order to be granted parallel observing time. Please refer to the HST Call for Proposals
for current policies and procedures concerning parallels.
Spatial scanning of stellar images upon the UVIS detector creates the potential for astrometry of unprecedented precision. Two representative scientific examples are parallax measurement of Cepheid variable stars (program 12679, Riess P.I.) and the astrometric wobble of a stellar binary (program 12909, Debes P.I.). Results from non-proprietary data of program 12679 (Riess, priv. comm.) indicate that differential astrometry a few times less precise than that set by diffraction and Poisson statistics are attainable. For HST
, a 2.4-m telescope, operating at 600 nm, the diffraction limit is Θ
/D = 51 mas. In the theoretical limit, astrometry in one dimension is approximately equal to the FHWM Θ
divided by the signal to noise ratio,
, where N is the number of photo-electrons recorded. If we adopt N equal to the full well of the UVIS CCD, ~64,000 e-
, times a trail of length 4000 pixels, i.e. N = 128 million e-
, then the theoretical astrometric limit is ~3 microarcsec per exposure. A more conservative estimate of ~13 microarcsec can be derived as follows: the nominal, state-of-the-art astrometric precision of a staring-mode exposure is ~0.01 pixel, so the astrometric precision of a 1000-pixel-long scan could be ~
or ~30 times smaller, which, for the 40 mas WFC3 UVIS pixels, is 13 microarcsec. In 2012 the TAC recommended programs 13101 and 12909, which anticipated a per-exposure precision of 30 to 40 microarcsec. We caution that considerable data analysis development is still ongoing from the scientific programs pioneering spatially-scanned WFC3 astrometry and observers must develop their own analysis software to reduce their images to obtain useful astrometric results. (See Riess et al. 2014
For those preparing a phase II program description, we recommend WFC3 ISR 2012-08.
Also, IR imaging with spatial scanning is discussed in Section 7.10.4
, and slitless spectroscopy with spatial scanning is discussed in Section 8.6
. See Figure 8.9
for a diagram provided in APT to assist observers planning spatial scan observations.
Note: starting in Cycle 24, the Exposure Time Calculator (ETC
) supports spatial scanning for UVIS and IR imaging and IR spectroscopy. (See WFC3 STAN issue 22
UVIS imaging has been shown to be highly effective in detecting faint point sources near bright point sources (WFC3 ISR 2011-03
). For a variety of narrow, medium, and wide filters, when a high signal-to-noise drizzled image of a star was scaled down by 10 magnitudes and shifted and added to the original image, the simulated faint companion could usually be seen for separations greater than 1.0 arcsec. Based on the annular signal-to-noise of the deep stellar image, 5 sigma detections of companions fainter by two magnitudes could be made at a separation of 0.1 arcsec. Theoretically, companions several magnitudes fainter could be detected at that separation in deeper images, but, in practice, variations in the PSF (point spread function) due to telescope breathing limit the detectability within about 0.3 arcsec of a bright star.
While Tiny Tim
modeling is available for the WFC3 UVIS detector, it has not been optimized to reproduce observed PSFs. See Section 6.6.4
for a discussion of on-going work to provide PSF models to observers.
The term DASH (for “drift-and-shift”, Momcheva et al., 2016
) has been adopted to describe the observing strategy of taking a series of WFC3/IR exposures of many targets within one orbit while the telescope is being guided under gyroscope control, thus avoiding the overhead cost of acquiring a new pair of guide stars for every slew between targets of greater than about 2 arcmin. A WFC3/IR sample sequence comprised of short exposure times is selected to avoid image smearing within each time step, and the differential samples in one exposure are later aligned and combined to compensate for the greater drift due to gyroscope control. The technique was designed to allow users to carry out shallow large-scale mosaic observations with the WFC3/IR camera, but has since been adapted to efficiently observe a collection of bright targets within a field ~1 deg across with WFC3/IR subarray apertures (which have shorter time steps for a given sample sequence) and WFC3/UVIS subarray apertures. More detailed discussions of HST guiding and considerations for program design are given in Section 7.10.6
. A brief discussion of WFC3/UVIS observations is given here.