An overview of the UVIS spectral elements was given in Section 2.3
. This section gives further details of the UVIS filters and grism. Table 6.2
contains a complete listing of the available spectral elements in the UVIS channel. Figures 6.3
show the effective throughput curves, including the filter transmission convolved with the OTA, WFC3 optics, and detector response. All of the UVIS filters are contained in a multi-wheel mechanism—identical to the mechanism on WFPC2—called the Selectable Optical Filter Assembly (SOFA). Values in Table 6.2
have been calculated for UVIS chip 2, which has a higher UV sensitivity, except in the cases of quad filters which are restricted to the A and B quadrants.
More detailed information on the throughput curves of all of the filters is given in Appendix A:WFC3 Filter Throughputs
; in particular, Section A.2.1
describes how to generate tabular versions of the throughput curves using synphot
. All measurements of the UVIS filters which involve wavelengths, as tabulated in Table 6.2
and plotted in Figures 6.3
and in Appendix A:WFC3 Filter Throughputs
, were done in air. The data have been converted to vacuum wavelengths using the formula given by D. C. Morton (1991, ApJS 77
, 119). It should also be noted that the laboratory measurements were done at a temperature of 20°
C, whereas the UVIS filters are operated on orbit at 0°
C. The temperature difference may lead to wavelength shifts that are no more than 0.14 nm in the worst cases, according to the filter manufacturing specifications.
The UVIS filters have been chosen to cover a wide variety of scientific
applications, ranging from color selection of distant galaxies to accurate photometry of stellar sources and narrow-band imaging of nebular gas. The set includes several very wide-band filters for extremely deep imaging, filters that match the most commonly used filters on WFPC2 and ACS (to provide continuity with previous observations), the SDSS filters, and filters that are optimized to provide maximum sensitivity to various stellar parameters (e.g., the Strömgren and Washington systems, and the F300X filter for high sensitivity to the stellar Balmer jump). There is a variety of narrow-band filters, which allow investigations of a range of physical conditions in the interstellar medium, nebulae, and solar system. A few of the narrow-band filters are also provided with slightly redshifted wavelengths, for use in extragalactic applications. Finally, there is a UV grism that provides slitless spectra with useful dispersion covering 200–400 nm (although the grism transmission spans the full wavelength range of the CCD).
Most of the UVIS filters, as well as the UVIS grism, are full-sized elements that
cover the entire UVIS field of view. However, in order to provide a larger number of bandpasses, there are five sets of “quad” filters, identified with “FQ” in the filter name, where each bandpass covers ~1/6 of the WFC3 UVIS field of view (i.e., each bandpass covers less than half of a single CCD chip). The quad filters are discussed in more detail below.
The UVIS channel is designed to be used with a single filter or grism in the light
path. Although possible in principle, unfiltered imaging, or imaging through a combination of two filters (from two different SOFA wheels), would lead to significantly degraded image quality and has not been tested; thus these options are not permitted. The F200LP filter provides a clear fused silica element that approximates unfiltered imaging.
While the red blocking in the WFC3 UV filters is generally very good, resulting in
negligible red leaks for hot objects (typically <<1% for targets with effective temperature Teff
> 10,000 K), the red leak can become significant in some filters for cooler targets (e.g., ~10% in F225W for a star with Teff
= 5000 K). More details are available in Section 6.5.2
; Table 6.5
in that section tabulates red-leak values as a function of stellar effective temperature.
Figure 6.3: Integrated system throughput of the WFC3 UVIS long-pass and extremely wide filters (top panel) and of the
wide-band filters covering 2000-6000 Å (bottom panel). The throughput calculations include the HST
OTA, WFC3 UVIS-channel internal throughput, filter transmittance, and the QE of the UVIS flight detector, and a correction factor to account for the gain sensitivity seen in SMOV4 on-orbit observations vs. TV3 ground tests. Throughputs in all plots below ~3200 Å contain contributions at the measured rate from all detected electrons, including UV multiple electrons.
As mentioned earlier, the WFC3 UVIS optics and CCDs have been optimized for
UV imaging. As such, the UV filters play a key role and considerable effort has been made to procure filters with the best possible characteristics, including maximum throughput, maximum out-of-band blocking, and minimal ghosts.
The UV filters include the shortest-wavelength F218W, intended for studies of the
ISM absorption feature; the wide F225W and F275W for broad-band UV imaging; the Strömgren u
(F336W) and Washington C (F390W) for stellar astrophysics; the extremely wide F300X for very deep imaging; and narrow bands such as F280N (Mg II) and the quad filters FQ232N and FQ243N (C II] and [Ne IV]).
The selection of extremely wide (X) and long-pass (LP) filters is suited for deep
imaging of faint sources. The ultra-wide F200LP filter is simply a fused-silica element with a UV-optimized anti-reflection coating which covers the UVIS channel’s entire spectral range (200-1000 nm). The F200LP filter is analogous to the clear filter on STIS.
WFC3’s maximum sensitivity to hot sources can be obtained by subtracting an
F350LP image from an F200LP, thereby creating a very broad ultraviolet bandpass. In Figure 6.7
, the blue curve shows the filter transmission for the F200LP filter, and the black curve shows the effective transmission for a F200LP minus F350LP difference image. For redder targets, some additional calibration may be necessary to account for differences in the transmission of the two filters longward of ~450 nm.
The F850LP filter is part of the Sloan Digital Sky Survey
filter set, and is the reddest of the ultra-wide filters.
The most commonly used WFPC2 and ACS wide filters are also included in the
WFC3 filter set. In addition to a wide V
-band filter (F606W), there is the Johnson-Cousins BVI
set (F438W, F555W, F814W).
The Sloan Digital Sky Survey
filter set (F475W, F625W, F775W, F850LP) is designed to provide high throughput for the wavelengths of interest and excellent rejection of out-of-band wavelengths. These filters provide wide, non-overlapping filter bands that cover the entire range of CCD sensitivity from blue to near-IR wavelengths.
The medium-band filters include the Strömgren set (u
, and y
), as well as some continuum bands needed for use with narrow-band imaging (F390M, FQ422M). The four 11% passband filters were added to the WFC3 UVIS set in order to cover the ~600-900 nm wavelength region with equal-energy filters. The “11%” refers to the filter bandwidths, which are ~11% of the central wavelength.
The WFC3 UVIS channel contains 36 different narrow-band filters, covering a
variety of species and most of the astrophysically interesting transitions, including Hα, Hβ, Hγ, He II, C II], [N II], [O I], [O II], [O III], [Ne IV], [Ne V], [S II], and Ca II. The methane absorption bands seen in planets, cool stars, and brown dwarfs are also covered.
Cosmological emission lines can be detected across a range of redshifts within the
bandpasses of the narrow-band filters. Table
lists the redshifts that can be probed using the specified narrow emission lines (hence, no entries for broad absorption bands or continuum or “off” bands). These redshift ranges are offered as a guide; exact values depend on the wavelengths of the filter cutoffs. Filter cutoffs used in Table
were defined using the passband rectangular widths (defined in Footnote d
of Table 6.2
). However, passband cutoffs were not centered on the filter pivot wavelengths λp
(defined in Section 9.3
), because red leaks shift the pivot wavelengths to longer wavelengths by 1-9% in some of the ultraviolet filters. Instead, the central wavelength for each filter was determined by maximizing the wavelength-integrated product of a rectangular passband of the specified width with the actual system throughput for the filter. In the most extreme case (FQ232N), the pivot wavelength of 241.3 nm is more than two bandpass widths to the red of the rectangular passband equivalent central wavelength (232.6 nm).
The WFC3 UVIS channel contains five quad filters: each is a 2×
2 mosaic of filter elements occupying a single filter slot, with each quadrant providing a different bandpass (typically narrow-band, although there are also several bandpasses intended for continuum measurements). The five quad filter sets on WFC3 significantly increase the number of available narrow-band filters. The WFC3 quad filters are listed in Table 6.4
with their readout amplifiers.
A quadrant nominally covers only 1/4 of the WFC3 total field of view or about
80", although edge effects (Figure 6.8
) result in an unvignetted field of about 1/6 of the field of view. The filter edges are out of focus on the focal plane, so light from multiple passbands reaches the detector in those areas.
In programs where targets are placed in different quadrants during a single orbit,
spacecraft maneuvers may be large enough to force a new guide star acquisition. Guide star acquisition overheads are described in Section 10.2
Figure 6.8: Quad filter edge effects (indicated by brackets). QUAD-FIX apertures
have reference points in the center of each quadrant.
Starting in Cycle 18, QUAD and 2K2-SUB apertures have had reference points in the center of the areas unaffected by filter edge effects
. QUAD-SUB apertures initially had quadrant-centered reference points, but will match the reference points in the QUAD apertures starting in cycle 20.
The UVIS channel has a UV grism (G280), the spare element from WF/PC-1. It
provides slitless spectra with a dispersion of about 1.4 nm/pix and a spectral resolution of about 70, over the 200-400 nm wavelength range, but with transmission in the zeroth order over the entire response of the CCD (see Figure 8.3
). Typically, a grism observation is accompanied by a direct image, for source identification and wavelength zero-point calibration; an ideal filter for the identification image is the F300X discussed above. Chapter 8
discusses WFC3 slitless spectroscopy in detail.
The design and manufacture of the UV filters was based on a careful balance of the
in- and out-of-band transmissions: in general, higher in-band transmission results in poorer suppression of out-of-band transmission, and vice versa. The WFC3 filters represent an attempt to achieve an optimum result, maximizing the in-band transmission while keeping the out-of-band transmission as low as possible in order to minimize red leaks.
below summarizes the red-leak levels for the WFC3 UV filters. The table lists the fraction of the total signal that is due to flux longward of 400 nm, as a function of effective temperature. This was calculated by convolving a blackbody of the given Teff
with the system throughput in the listed filter. As can be seen from the table, red leaks should not be an issue for observations of any objects taken with F275W or F336W. The other UV filters have some red leaks, whose importance depends on stellar temperature. The red leaks in F218W and F300X, for example, exceed ~1% for objects cooler than ~6000 K, while in F225W the red leak reaches ~1% for objects with even cooler temperatures. The most extreme red leaks arise from F218W and F225W observations of objects with Teff
of ~4000 K or cooler, necessitating appropriate corrections.
The WFC3 UVIS channel exhibits three different types of optical ghosts: a) those
due to reflections between the CCD front surface and the two detector package windows; b) those due to reflections between the window surfaces; and c) those due to reflections within the particular filter in use.
When a point source is positioned in the lower right quadrant of the UVIS detector,
out-of-focus reflections between the CCD and windows appear along a diagonal from the source towards the upper left, well removed from the source. These figure-eight shaped ghosts gradually move outside the field of view as the target moves out of the lower right corner. They contain a few percent of the flux of the target. The ghosts of several bright stars are visible in the exposure of 47 Tuc shown in Figure 6.9
To prevent the worst effects of window ghosts, avoid placing very bright targets on
the D quadrant. Also, pay attention to the location of key science targets if bright sources are in the lower right area of the field of view. The production of window ghosts has been modeled and an aid to observers has been produced to enable them to estimate the position of ghosts for a given field of view and ORIENT (WFC3 ISR 2011-16
). If necessary, ORIENT special requirements can be imposed within APT at the Phase II proposal preparation stage to control the positioning of bright sources on the detector.
Filter ghosts for the WFC3 filters were specified to be less than 0.2%, and in most
cases were measured during ground testing to be less than 0.1%. A few filters however, were found during ground testing to have ghosts that exceeded the specification. Some of these, the ones deemed highest priority, were remanufactured and installed in the SOFA. Consequently, there are a relatively small number of filters that may be of concern for ghosts. These are listed in Table 6.6
. They have been retained in the flight instrument either because they were of lower scientific priority, or because the ghost level was deemed acceptable in light of the otherwise excellent performance characteristics of the filters (e.g., in- and out-of-band transmission, sharpness of bandpass edges). While some scientific programs (e.g., stellar photometry) may be unaffected by filter ghosts, others (e.g., observations of extended targets or faint objects adjacent to bright ones) could be adversely affected. In such cases, extra planning and/or data-analysis efforts may be needed, e.g., combining images taken at different dither positions and/or roll angles, or applying a deconvolution algorithm.