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| Wide Field Camera 3 Instrument Handbook for Cycle 17 | ||||
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6.5 UVIS Spectral Elements
6.5.1 Filter and Grism Summary
An overview of the UVIS spectral elements, and of the process by which they were selected, 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 6.2 through 6.5 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 package called the Selectable Optical Filter Assembly (SOFA).
More detailed information on the throughput curves of all of the filters is given in Appendix A. All measurements of the UVIS filters which involve wavelengths, as tabulated in Table 6.2 and plotted in Figures 6.2 through 6.5 and in Appendix A, were done in air. In evaluating critical cases (e.g., narrow-band imaging of redshifted galaxies), one should convert the air wavelengths to vacuum using, for example, 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 20ºC, whereas the UVIS filters will be operated on orbit at 0ºC; this may lead to wavelength shifts which are expected to be small.
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 will 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 covering 200-400 nm.
Note that, in contrast to ACS, WFC3 does not have any polarizers or ramp filters.
Table 6.2: WFC3/UVIS Filters and Grism.
Name1
(Å)
TransmissionF200LP F300X F350LP F475X F600LP F850LP F218W F225W F275W F336W F390W F438W F475W F555W F606W F625W F775W F814W F390M F410M FQ422M F467M F547M F621M F689M F763M F845M FQ232N FQ243N F280N F343N F373N FQ378N FQ387N F395N FQ436N FQ437N F469N F487N FQ492N F502N FQ508N FQ575N FQ619N F631N FQ634N F645N F656N F657N F658N F665N FQ672N F673N FQ674N F680N FQ727N FQ750N FQ889N FQ906N FQ924N FQ937N F953N G280
1The spectral-element naming convention is as follows for both the UVIS and IR channels. All filter names begin with F, and grisms with G; if the filter is part of a four-element quad mosaic, a Q follows F. Then there is a three-digit number giving the nominal effective wavelength of the bandpass, in nm (UVIS channel) or nm/10 (IR channel). (For long-pass filters, the number is instead the nominal blue cut-off wavelength in nm.) Finally, for the filters, one or two letters indicate the bandpass width: X (extremely wide), LP (long pass), W (wide), M (medium), or N (narrow).
2Filters intended for imaging in a red-shifted bandpass are given descriptions similar to the following: "z (H+ [N II])".
3"Pivot wavelength" is a measure of the effective wavelength of a filter (see A. Tokunaga & W. Vacca 2006, PASP, 117, 421). It is calculated here based only on the filter transmission. Values are approximate for the long-pass filters. All wavelength measurements in this table were made in air.
4Full width at 50% of peak transmission for wide and medium bands, and at 10% of peak transmission for narrow bands. For long-pass filters, the widths are approximate and include convolution with the detector QE.
5km-agt (km-amagat) is a unit of vertical column density, equal to 2.69×1024 molecules/cm2.
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 "quad" filters, identified with "FQ" in the filter name, where each bandpass covers ~1/4 of the WFC3 UVIS field of view (i.e., each bandpass covers 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.
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 Filter Red Leaks; Table 6.3 in that section tabulates red-leak values as a function of stellar effective temperature.
Figure 6.2: Integrated system throughput of the WFC3 UVIS long-pass and extremely wide filters (top panel) and of the wide-band filters covering 2000-6000 A (bottom panel). The throughput calculations include the HST OTA, WFC3 UVIS-channel internal throughput, filter transmittance, and the QE of the UVIS flight detector. Throughputs in all plots below ~3200 A include correction for quantum yield. Instrument throughput is not yet well-characterized below 2000 A and above 10,000 A.
Figure 6.3: Integrated system throughput of the WFC3 UVIS wide-band filters covering 4000-10,000 A (top panel) and the medium-band filters (bottom panel). The throughput calculations include the HST OTA, WFC3 UVIS-channel internal throughput, filter transmittance, and the QE of the UVIS flight detector.
Figure 6.4: Integrated system throughput of the WFC3 UVIS narrow-band filters covering 2000-4500 A (top panel) and the narrow-band filters covering 4500-6000 A (bottom panel). The throughput calculations include the HST OTA, WFC3 UVIS-channel internal throughput, filter transmittance, and the QE of the UVIS flight detector.
Figure 6.5: Integrated system throughput of the WFC3 UVIS narrow-band filters covering 6000-6800 A (top panel) and the narrow-band filters covering 6600-9600 A (bottom panel). The throughput calculations include the HST OTA, WFC3 UVIS-channel internal throughput, filter transmittance, and the QE of the UVIS flight detector.
UV Filters
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]).
There is also an ultra-wide F200LP filter; it 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 6.6, the gray 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.
Figure 6.6: Sensitivity of F200LP-F350LP compared to other UV filters (F225W, F275W and F300X). Dashed curves show blackbody functions for 20,000 and 50,000 K.
Wide-band 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 (SDSS) griz 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 will provide wide, non-overlapping filter bands that cover the entire range of CCD sensitivity from blue to near-IR wavelengths.
Medium-band Filters
The medium-band filters include the Strömgren set (u, v, b, 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.
Narrow-band Filters
The WFC3 UVIS channel contains 36 different narrow-band filters, covering a variety of species and most of the astrophysically interesting transitions, including H
, H
, 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.
Quad Filters
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).
A quadrant nominally covers only 1/4 of the WFC3 total field of view or about 80"×80", although edge effects will result in an unvignetted field somewhat smaller than 1/4 of the field of view. The five quad filters on WFC3 significantly increase the number of available narrow-band filters. The WFC3 quad filters are identified by the prefix "FQ" in the filter name in Table 6.2.
Grism
The UVIS channel has a UV grism (G280), a spare 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. Typically, a grism observation will be accompanied by a direct image, for source identification and wavelength calibration; an ideal filter for the identification image would be the F200LP discussed above. Chapter 8 discusses WFC3 slitless spectroscopy in detail.
6.5.2 Filter Red Leaks
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.
Table 6.3 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 stellar effective temperature. As can be seen from the table, red leaks should not be an issue for observations of any target taken with F275W or F336W. The other UV filters will have some red leak, whose importance depends on stellar temperature. The red leak in F218W, for example, will exceed ~1% for objects cooler than ~7000 K, while in F300X the red leak reaches ~1% for objects as hot as 10,000 K. The most extreme red leaks will arise from F218W and F225W observations of objects with Teff of ~5000 K or cooler, necessitating appropriate corrections.
Table 6.3: Fraction of flux longward of 400 nm as a function of effective temperature.
6.5.3 Ghosts
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 between the surfaces of the particular filter in use.
Window Ghosts
When a point source is positioned at the lower right of the field of view, reflections between the CCD and windows appear along the diagonal from the center of the field of view towards the upper left; these ghosts gradually move outside the field of view as the target source moves out of the lower right corner. Smaller ghosts appear closer to the target source: they are due to reflections between the window surfaces (see Figure 6.10 for an image showing these ghosts).
Filter Ghosts
Filter ghosts, however, can occasionally fail to satisfy the specification. During initial instrument-level ground tests, while most of the WFC3 UVIS filters performed consistently with, or exceeding, expectations, a few of the filters were found to have filter ghosts. The highest-priority filters affected by ghosts were re-manufactured and re-installed into the SOFA.
The levels of filter ghosts in the majority of the final flight filter set are now <0.1%, well within the specified requirements. However, a subset of the final flight filters still have ghosts exceeding the specification, which calls for <0.2% of the total flux to reside in the ghost. These filters are listed in Table 6.4. 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.
Table 6.4: Filters exceeding the filter ghost requirement, measured during ground testing of the integrated instrument (see WFC3 ISR 2007-09).
1Laboratory measurement of stand-alone filter.
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Space Telescope Science Institute http://www.stsci.edu Voice: (410) 338-1082 help@stsci.edu |
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