STScI

Advanced Camera for Surveys Instrument Handbook for Cycle 14

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Chapter 5:
Polarimetry, Coronagraphy and Prism/Grism Spectroscopy

5.1 Polarimetry
5.2 Coronagraphy
    5.2.1 Coronagraph Design
    5.2.2 Acquisition Procedure and Pointing Accuracy
    5.2.3 Vignetting and Flat Fields
    5.2.4 Coronagraphic Performance
    5.2.5 Residual Light Subtraction
    5.2.6 The Off-Spot PSF
    5.2.7 Occulting Spot Motions
    5.2.8 Planning ACS Coronagraphic Observations
    5.2.9 Choice of Filters for Coronagraphic Observations
5.3 Grism/Prism Spectroscopy
    5.3.1 WFC G800L
    5.3.2 HRC G800L
    5.3.3 HRC PR200L
    5.3.4 SBC PR110L
    5.3.5 SBC PR130L
    5.3.6 Observation Strategy
    5.3.7 Extraction and Calibration of Spectra

In this chapter we provide an overview of the special observing capabilities offered by ACS. These capabilities are optical and near-UV imaging polarimetry, coronagraphy with an aberrated beam coronagraph and low resolution (R~100) optical and near-UV spectroscopy.

5.1 Polarimetry

The Advanced Camera has a straightforward imaging polarimetric capability. Polarization observations require a minimum of three images taken using polarizing optics with different polarization characteristics in order to solve for the source polarization unknowns (polarization degree, position angle and total intensity). To do this, ACS offers two sets of polarizers, one optimized for the blue (POLUV) and the other for the red (POLV). These polarizers can be used in combination with most of the ACS filters (see Table 5.2) allowing polarization data to be obtained in both the continuum and in line emission; and to perform rudimentary spectropolarimetry by using the polarizers in conjunction with the dispersing elements. (Due to the large number of possibilities in combination with ramp and dispersing elements, and heavy calibration overheads, observers wishing to use those modes should request additional calibration observations). For normal imaging polarization observations, the target remains essentially at rest on the detector with a suitable filter in beam, and an image is obtained with each of the appropriate polarizing elements in turn. The intensity changes between the resulting images provide the polarization information.

Each set of polarizers comprises three individual polarizing filters with relative position angles 0º, 60º and 120º. The polarizers are designed as aplanatic optical elements and are coated with "Polacoat 105UV" for the blue optimized set and HN32 polaroid for the red set. The blue/near-UV optimized set is also effective all through the visible region, giving a useful operational range from approximately 2000Å to 8500Å. The second set is optimized for the visible region of the spectrum and is fully effective from 4500Å to about 7500Å.

The relative performance of the UV-optimized versus the visible optimized polarizers is shown in Figure 5.1. The visible polarizers clearly provide superior rejection for science in the 4500-7500Å bandpass, while the UV optimized polarizers deliver lower overall rejection across a wider range into the near-UV, 2000-7500Å. While performance of the polarizers begins to degrade at wavelengths longer than about 7500Å, useful observations should still be achievable to approximately 8500Å in the red. In this case, allowance for imperfect rejection of orthogonally polarized light should be made at the analysis stage.

A further caveat is that imperfections in the flat fields of the POLVIS polarizer set have been found which may limit the optimal field of view somewhat. Potential users are encouraged to check the STScI ACS web site for the latest information.

To first approximation, the ACS polarizers can be treated as three essentially perfect polarizers. The Stokes parameters (I, Q, U) in the most straightforward case of three images obtained with three perfect polarizers at 60º relative orientation, can be computed using simple arithmetic.

Using im1, im2, and im3 to represent the images taken through the polarizers POL0, POL60, and POL120 respectively, the Stokes parameters are as follows:





These values can be converted to the degree of polarization P and the polarization angle , measured counterclockwise from the x axis as follows:

A more detailed analysis, including allowance for imperfections in the polarizers may be found in Sparks & Axon 1999 PASP, 111, 1298. They find that the important parameter in experiment design is the product of expected polarization degree and signal-to-noise. A good approximation for the case of three perfect polarizers oriented at the optimal 60º relative position angles (as in ACS) is that the error on the polarization degree P (which lies in the range 0 for unpolarized to 1 for fully polarized) is just the inverse of the signal-to-noise per image. Specifically, they found

where is the signal to noise of the ith image; and

The above discussion is for ideal polarizers with no instrumental polarization. Of course, the reality is that the polarizer filters, especially the UV polarizer, has significant leakage of cross-polarized light. And the instrumental polarization of the HRC ranges from a minimum of 4% in the red to 14% in the far-UV, while that of the WFC is ~2% (see ISR ACS 2004-09). Other effects, such as phase retardance in the mirrors, may be significant as well. Please consult the ACS Data Handbook, STScI web pages, and ISRs for more detailed information.

Figure 5.1: Throughput and rejection of the ACS Polarizers. In the top two boxes, the upper curve is the parallel transmission, while the lower curve is the perpendicular transmission. The bottom panel shows the logarithm of the ratio of perpendicular to parallel transmission


 

The implementation of the ACS polarizers is designed for ease of use. The observer merely selects the camera (either HRC or WFC), the spectral filter, and then takes images stepping through the three filters of either the either the VIS set (POL0V, POL60V, POL120V) or the UV set (POL0UV, POL60UV, POL120UV). Once the camera and polarizer are specified, the scheduling system automatically generates slews to place the target in the optimal region of the field of view.

Since the ACS near-UV and visible filter complement is split between two filter wheels, there are restrictions on which filters the polarizer sets can be combined with. The choices available were determined by the relative performance of the polarizers and the near-UV limitations of the WFC resulting from the silver mirror coatings.

The near-UV optimized polarizers are mounted on Filter Wheel 1 and may be crossed with the near-UV filter complement, which are mounted on Filter Wheel 2. The visible optimized polarizers are mounted on Filter Wheel 2 and can be crossed with filters on Filter Wheel 1, namely the primary broadband filters, and discrete narrowband filters H, [OII] and their continuum filters. Due to the calibration overhead required, it is not planned to support the use of ramp filters with either polarizer set. GOs are, therefore, required to include calibration observations, if they plan to use the ramp filters with the polarizers.

The polarizer sets are designed for use on the HRC where they offer a full unvignetted field of view, 29×26 arcsec with any of the allowable filter and coronagraph combinations including those ramps and spectroscopic elements that may also be used on the HRC (although see above re. additional calibrations).

The same allowable combinations, either UV or visible optimized, may also be used on the WFC where an unvignetted field of view of diameter 70 arcsec is obtained. This does not fill the field of view of the WFC due to the small size of the polarizing filters. However it does offer an area approximately five times larger than that obtained on the HRC. In order to avoid the gap between the WFC CCDs, and to optimize the readout noise and CTE effects, the scheduling system will automatically slew the target to roughly pixel (3096,1024) on the WFC1 CCD whenever the WFC aperture is selected in conjunction with one of the polarizers. Also, to reduce camera overhead times, only a 2048x2048 subimage centered on the target location will be readout from WFC1 (see Table 5.1).

Occasionally observers will ask to obtain non-polarized images at the same physical location on the detector as their polarized images. This is completely straight forward for the HRC; one merely takes the exposure without the polarizer filter. However, for the WFC it is more complicated because specifying WFC together with a polarizer automatically invokes a large slew, whereas no slew is performed when the polarizer filter is omitted. To obtain a non-polarizer image at the same physical detector location as the polarizer image in the WFC, one merely needs to specify the aperture as WFC1-2K instead of WFC (see Table 5.1).


Table 5.1: Examples of Polarizer and non-Polarizer Exposures in a Phase II Proposal
Aperture Filters Comment
HRC F606W, POL0V 1024x1024 image centered at usual HRC aperture.
HRC F606W, POL60V Same but with POL60V.
HRC F606W, POL120V Same but with POL120V.
HRC F606W Non-polarizer image centered at same detector location as polarizer exposure.
WFC F606W, POL0V Target automatically placed at WFC1 pixel (3096,1024); 2048x2048 image.
WFC F606W, POL60V Same but with POL60V.
WFC F606W, POL120V Same but with POL120V.
WFC1-2K F606W Non-polarizer image at same detector location. Target at WFC1 pixel (3096,1024); 2048x2048 image.


Table 5.2: Filters that can be used in conjunction with the ACS Polarizers
Polarizer set
Filters
Filter Comments
POL0UV POL60UV POL120UV F220W F250W F330W F435W F814W HRC NUV short HRC NUV long HRC U Johnson B broad I
POL0V POL60V POL120V F475W F606W F625W F658N F775W SDSS g Johnson V SDSS r H SDSS i

The filters specified in Table 5.2 are those that we expect users to choose for their polarization observations. We will calibrate the most popular of these filters. Filter combinations not on this list will most probably not be calibrated, so potential users who have a strong need for such a polarizer/filter combination should include any necessary calibrations themselves.

We anticipate that the most accurate polarization observations will be obtained in the visible band (i.e., F606W) with the HRC and the visible polarizers. This mode has the advantages of a very high rejection of perpendicular polarization, and known mirror coatings with readily modeled properties. The WFC may be capable of similar accuracy to the HRC, however its proprietary mirror coatings will make modeling of the polarization properties, and hence calibration, much more difficult (e.g., unknown phase retardance effects in the WFC IM3 mirror are a concern).

Polarimetry in the UV will be more challenging for a number of reasons. The UV polarizer has relatively poor rejection in the UV, and the instrumental polarization of the HRC, which is 4% - 7% in the visible, rises to 8% - 9% in the UV, and reaches 14% at 2200 Å (see ACS ISR 04-09). Far-UV polarimetry will be especially challenging since the polarizer properties were not well-characterized shortwards of 2800 Å, and appear to change rapidly with wavelength. Moreover, the low UV transmission of the UV polarizer, and the poor polarization rejection in the far-red, work together to exacerbate redleaks which are normally seen in the far-UV spectral filters.

The polarizer filters contribute a weak geometric distortion which rises to about 0.3 pixel near the edges of the HRC field-of-view. This is caused by a weak positive lens in the polarizers, which is needed to maintain proper focus when multiple filters are in the beam. In addition, the visible polarizer has a weak ripple structure which is related to manufacture of its polaroid material; this contributes an additional +/-0.3 pixel distortion with a very complex structure (see ACS ISR 04-10 and ACS ISR 04-11). All these geometric effects are correctable with the drizzle software, however, astrometry will likely be less accurate in the polarizers due to residual errors and imperfect corrections.

5.2 Coronagraphy

The ACS High Resolution Camera (HRC) has a user-selectable coronagraphic mode for the imaging of faint objects (circumstellar disks, substellar companions) near bright point sources (stars or luminous quasar nuclei). The coronagraph suppresses the diffracted light (diffraction spikes and rings) of the central source to below the level of the scattered light, most of which is caused by surface errors in the HST optics. The coronagraph was added after ACS construction began, at which point it was impossible to insert it into the aberration-corrected beam. Instead, the system is used in the aberrated beam, which is corrected after the coronagraph. While not as efficient as a corrected-beam coronagraph, especially for imaging close to the central source, it does provide a significant improvement to the high-contrast imaging capabilities of HST. Care must be taken, however, to design an observation plan that properly optimizes the coronagraph's capabilities and accounts for its limitations.

Figure 5.2: Schematic layout of the ACS HRC coronagraph. The upper left inset shows a schematic of the coronagraph mechanism that can be flipped in-and-out of the HRC optical path.


 

5.2.1 Coronagraph Design

A schematic layout of the ACS coronagraph is shown in Figure 5.2. The aberrated beam from the telescope first encounters one of two occulting spots. The beam continues to the M1 mirror, which forms an image of the HST entrance pupil on the M2 mirror, which corrects for the spherical aberration in the HST primary mirror. The coronagraph's Lyot stop is placed in front of M2. A fold mirror directs the beam onto the CCD detector. The field is 29" by 26" with a mean scale of 0.026"/pixel (geometric distortion results in effectively non-square pixels). The coronagraph can be used over the entire HRC wavelength range of =2000-10,000Å using a variety of broad-to-narrowband filters.

The occulting spots are placed in the plane of the circle of least confusion, near where the unaberrated HST focal plane would be. At this location the balance of defocus and spherical aberration provides a good compromise between maximal occulted flux and minimal spot radius. The angular extent of the PSF in this plane necessitates larger spots than would be used in an unaberrated system (Figure 5.3). The ACS spots are solid (unapodized) metallic coatings deposited on a glass substrate (which reduces throughput by 4.5%). The smaller spot is 1.8" in diameter and is at the center of the field. It is selected with the aperture CORON-1.8. A 3.0" diameter spot is near a corner (Figure 5.4) and is designated CORON-3.0. The smaller spot is used for the majority of the coronagraphic observations, as it allows imaging closer to the central source. The larger one may be used for very deep imaging of bright targets with less saturation around the spot edge than would occur with the smaller spot. Its position at the edge of the field also allows imaging of material out to 20" from the central source. The Lyot stop is located just in front of the M2 aberration correction mirror, where an image of the HST primary is formed. The stop is a thin metal mask that covers all of the diffracting edges in the HST system at the reimaged pupil (outer aperture, secondary mirror baffle, secondary mirror support vanes, and primary mirror support pads). The sizes of the stop and occulting spots were chosen to reduce the diffracted light below the level of the scattered light, which is unaltered by the coronagraph. The stop reduces the throughput by 48%, and it broadens the field PSF due to the smaller aperture and larger central obscuration relative to the beam diameter. The spots and Lyot stop are located on a panel attached to the ACS calibration door mechanism, which allows them to be flipped out of the beam when not in use. The inside surface of this door can be illuminated by a lamp to provide flat field calibration images for direct-mode imaging. However, the arrangement prevents the acquisition of internal flat fields in coronagraphic mode.

Figure 5.3: Computed point spread functions at the plane of the occulting spots through filters F435W and F814W. The elliptical, cross-shaped patterns in the cores are due to astigmatism at the off-axis location of the ACS aperture. It is corrected later by the ACS optics. The sizes of the two occulting spots (D=1.8" and 3.0") are indicated. Logarithmic intensity scaled.


 

In addition to the combination of the occulting spots and Lyot stop that comprise the coronagraph, there is a 0.8" wide, 5" long occulting finger (OCCULT-0.8) permanently located at the window of the CCD dewar. It does not provide any suppression of diffracted light because it occurs after the Lyot stop. It was to be used to image closer to stars than is possible with the occulting spots while preventing saturation of the detector. However, because the finger is located some distance from the image plane, there is significant vignetting around its edges, reducing its effectiveness. Originally aligned to cover the central portion of the 3.0" spot, shifting of the spots relative to the beam during launch now places the finger along that spot's edge. Because of vignetting and the sensitivity of the PSF wings to the centering of the star, unocculted, saturated observations of sources will likely be more effective than using the occulting finger.

Figure 5.4: Region of the Orion Nebula observed with the coronagraph in filter F606W. The silhouettes of the occulters can be seen superposed against the background nebulosity. The 1.8" spot is located at the center and the 3.0" spot towards the top. The finger is aligned along one edge of the larger spot. This image has not been geometrically corrected, so the spots appear elliptical.


 

5.2.2 Acquisition Procedure and Pointing Accuracy

The central source must be placed precisely behind the occulting spot to ensure the proper suppression of the diffracted light. The absolute pointing accuracy of HST is about 1", too crude to ensure accurate positioning. An on-board acquisition procedure, borrowed from the STIS coronagraph, is used to provide better alignment. The observer must request an acquisition image immediately before using the coronagraph and must specify a combination of filter and exposure time that provides an unsaturated image of the source. An acquisition image is taken by specifying HRC-ACQ as the aperture and ACQ as the opmode in APT.

Beginning in Cycle 12, acquisition images are taken with the coronagraphic masks inserted. The star is imaged within a predefined 200 x 200 pixel (5"x5") subarray near the small occulting spot. Two identical exposures are taken, each of the length specified by the observer (rather than each being half the length specified, as they would be for a conventional CR-SPLIT). From these two images, the on-board computer selects the minimum value for each pixel as a crude way of rejecting cosmic rays. The result is then smoothed with a 3x3 pixel box and the maximum pixel in the subarray is identified. The center-of-mass centroid is then computed for the unsmoothed image within a 5x5 pixel box centered on this pixel. Based on this position, the telescope is then slewed to place the star behind the occulting spot.

Because the coronagraphic masks are in place during acquisition, throughput is decreased by 52% relative to a non-coronagraphic observation. Also, the PSF is broader than in the normal imaging mode due to the larger obscurations in the Lyot stop, resulting in a lower relative peak pixel value (see Section 5.2.6). Care must be taken to select a combination of exposure time and filter that will prevent saturation of the star while providing enough flux to provide a good centroid measurement. A peak pixel flux of 2000 e- should be considered the minimum while 50K e- is a safe maximum (the HRC saturation limit is ~140K e-). Narrowband filters can be used, but for the brightest targets crossed filters are required. Allowable filter combinations for acquisitions are F220W+F606W, F220W+F550M, and F220W+F502N, in order of decreasing throughput. Be warned that the calibration of these filter combinations is poor, so estimated count rates from SYNPHOT or the APT ETC should be considered to be a factor of two off (either high or low).

Initial results from multiple on-orbit observations indicate that the combined acquisition and slew errors are on the order of ±0.25 pixels (±6 mas). While small, these shifts necessitate the use of subpixel registration techniques to subtract one coronagraphic PSF from another (Section 5.2.5). The position of the spots relative to the detector also varies over time. This further alters the PSF, resulting in subtraction residuals.

5.2.3 Vignetting and Flat Fields

ACS coronagraphic flat fields are different from the standard flats due to the presence of the occulting spots and alteration of the field vignetting caused by the Lyot stop. The large angular extent of the aberrated PSF results in vignetting beyond one arcsecond of the spot edge (Figure 5.4), which can be corrected by dividing the image by the spot pattern (Figure 5.5). Unfortunately, the spots move over time relative to the other flat field features (see Section 5.2.7). To account for this, separate flat fields have been derived that contain just the spot patterns (spot flats) and the remaining features (P-flats). For a full discussion see ACS ISR 04-16.

The ACS pipeline will divide images by the P-flat. P-flats specific to the coronagraph have been derived from either ground-based or on-orbit data for filters F330W, F435W, F475W, F606W, and F814W. Other filters use the normal flats, which may cause some small-scale errors around dust diffraction patterns. As of the time of writing, the pipeline does not divide the images by spot flats, which must be manually applied by the observer as described in the ACS Data Handbook. Spot flats for the filters listed above are available for download from the ACS reference files web page. For other filters, the available spot flat closest in wavelength should be used. The spot flat must be shifted by an amount listed on the reference files page to account for motions of the occulting spots.

Because coronagraphic P-flats and spot flats exist only for the few filters listed above, observers are encouraged to use those filters. It is unlikely that flat fields for other filters will be available in the future.

Figure 5.5: (Left) Region of the Orion Nebula around the D=1.8" spot. The spot edge appears blurred due to vignetting. The image has not been geometrically corrected. (Right) The same region after the image has been corrected by dividing the flat field. The interior of the spot has been masked.


 

5.2.4 Coronagraphic Performance

Early in Cycle 11, coronagraphic performance verification images were taken of the V=0 star Arcturus (Figures 5.6 & 5.7). This star has an angular diameter of 25 mas and is thus unresolved by the coronagraph. The coronagraphic image of a star is quite unusual. Rather than appearing as a dark hole surrounded by residual light, as would be the case in an aberration-free coronagraph, the interior of the spot is filled with a diminished and somewhat distorted image of the central source. This is due to correction by the M2 mirror of aberrated light from the star that is not blocked by the spot. The small spot is filled with light, while the large one is relatively dark. Broad, ring-like structures surround the spots, extending their apparent radii by about 0.5". These are due to diffraction in the wings of the aberrated PSF by the occulting spot itself. A consequence of these features is that stars may saturate the interior and edges of the spot within a short time. Within the small spot, the brightest pixels can become saturated in less than one second for a V=0.0 star, while pixels at edge of the larger spot will saturate in about 14 seconds.

Figure 5.6: Geometrically corrected (29" across) image of Arcturus observed in F814W behind the 1.8" spot. This is a composite of short, medium, and long (280s) exposures. The "bar" can be seen extending from the upper left to lower right. The shadows of the occulting finger and large spot can be seen against the scattered light background. Logarithmic intensity scale.


 

The measured radial surface brightness profiles (Figure 5.8) show that the coronagraph is well aligned and operating as expected. The light diffracted by the HST obscurations is suppressed below the level of the scattered light - there are no prominent diffraction spikes, rings, or ghosts beyond the immediate proximity of the spots. At longer wavelengths (>600 nm) the diffraction spikes appear about as bright as the residual scattered light (at longer wavelengths, the diffraction pattern is larger and therefore not as well suppressed by the coronagraph). The spikes are more prominent in images with the large spot than the small one. This can be explained by the fact that the Lyot stop is not located exactly in the pupil plane but is instead slightly ahead of it, so the beam can "walk" around the stop depending on the field angle of the object. Because the large spot is at the edge of the field, the beam is slightly shifted, allowing more diffracted light to pass around the mask edges.

The residual background is dominated by radial streaks that are caused primarily by scattering from zonal surface errors in the HST mirrors. This halo increases in brightness and decreases in size towards shorter wavelengths. One unexpected feature is a diagonal streak or "bar" seen in both direct and occulted star images. It is about 5x brighter than the mean azimuthal surface brightness in the coronagraphic images. This structure was not seen in the ground-test images and is likely due to scattering introduced by the HST optics. There appears to be a corresponding feature in STIS as well.

Figure 5.7: Regions around the occulting spots in different filters. The occulting finger can be seen in the 3" spot images. Logarithmic intensity scaled.


 
Figure 5.8: Surface brightness plots derived by computing the median value at each radius. The brightness units are relative to the total flux of the star. The direct profile is predicted; the coronagraphic profiles are measured from on-orbit images of Arcturus. "Coronagraph-star" shows the absolute median residual level from the subtraction of images of the same star observed in separate visits.


 

Figure 5.8 (continued)

5.2.5 Residual Light Subtraction

While the coronagraph suppresses the diffracted light from the central star, the scattered light still overwhelms faint, nearby sources. It is possible to subtract most of the remaining halo using an image of another occulted star. PSF subtraction has been successfully used with images taken by other HST cameras, with and without a coronagraph. The quality of the subtraction depends critically on how well the target and reference PSFs match.

As mentioned above, for any pair of target and reference PSF observations there is likely to be a difference of 5-20 mas between the positions of the stars. Because the scattered light background is largely insensitive to small errors in star-to-spot alignment (it is produced before the coronagraph), most of it can be subtracted if the two stars are precisely registered and normalized. Due to the numerous sharp, thin streaks that form the scattered light background, subtraction quality is visually sensitive to registration errors as small as 0.03 pixels (0.75 mas). To achieve this level of accuracy, the reference PSF may be iteratively shifted and subtracted from the target until an offset is found where the streaks are minimized. This method relies on the judgment of the observer, as any circumstellar material could unexpectedly bias a registration optimization algorithm. A higher-order sampling method, such as cubic convolution interpolation, should be used to shift the reference PSF by subpixel amounts; simpler schemes such as bilinear interpolation degrade the fine PSF structure too much to provide good subtractions.

Normalization errors as small as 1-4% between the target and reference stars may also create significant residuals. However, derivation of the normalization factors from direct photometry is often not possible. Bright, unocculted stars will be saturated in medium or broadband filters at the shortest exposure time (0.1 sec). An indirect method uses the ratio of saturated pixels in unocculted images (the accuracy will improve with greater numbers of saturated pixels). A last-ditch effort would rely on the judgment of the observer to iteratively subtract the PSFs while varying the normalization factor.

In addition to registration offsets, positional differences can alter the diffraction patterns near the spots' edges. The shape and intensity of these rings are very sensitive to the location of the star relative to the spot. They cannot be subtracted by simply adjusting the registration or normalization. These errors are especially frustrating because they increase the diameter of the central region where the data are unreliable. The only solution to this problem is to observe the target and reference PSF star in adjacent orbits without flipping the masks out of the beam between objects.

Color differences between the target and reference PSF can be controlled by choosing an appropriate reference star. As wavelength increases, the speckles that make up the streaks in the halo move away from the center while their intensity decreases (Figure 5.7). The diffraction rings near the spots' edges will expand as well. These effects can be seen in images through wideband filters - a red star will appear to have a slightly larger PSF than a blue one. Thus, an M-type star should be subtracted using a similarly red star - an A-type would result in significant residuals. Even the small color difference between A0V and B8V stars, for example, may be enough to introduce bothersome errors (Figure 5.9).

A focus change can also alter the distribution of light in the PSF. The telescope focus changes over time scales of minutes to months. Within an orbit, the separation between the primary and secondary mirrors varies on average by 3 µm (resulting in 1/28 wave RMS of defocus @ =0.5 µm) - an effect called breathing. This is caused by the occultation of the telescope's field of view by the warm Earth, which typically occurs during half of each 96-minute orbit. This heats HST's interior structure, which expands. After occultation the telescope gradually shrinks. Large changes in the pointing attitude relative to the Sun can also introduce 3-10 µm of expansion, which decays back to normal over several orbits. The main result of these small focus changes is the redistribution of light in the wings (Figure 5.10).

Figure 5.9: Predicted absolute mean subtraction residual levels for cases where the target and reference stars have color mismatches. The brightness units are relative to the total flux of the target star.


 
Figure 5.10: Predicted absolute mean subtraction residual levels for cases where the target and reference stars are imaged at different breathing-induced focus positions. The offset (0.75 or 2.5 µm) from perfect focus (0 µm) is indicated with respect to the change in the primary-secondary mirror separation (the typical breathing amplitude is 3-4 µm within an orbit). The brightness units are relative to the total flux of the target star.


 

Plots of the azimuthal median radial profiles after PSF subtraction are shown in Figure 5.8. In these cases, images of Arcturus were subtracted from others of itself taken a day later. The images were registered as previously described. Combined with PSF subtraction, the coronagraph reduces the median background level by 250-2500x, depending on the radius and filter. An example of a PSF subtraction is shown in Figure 5.11. The mean of the residuals is not zero. Because of PSF mismatches, one image will typically be slightly brighter than the other over a portion of the field (such as shown in Figure 5.12). The pixel-to-pixel residuals can be more than 10x greater than the median level (Figure 5.13). Note that these profiles would be worse if there were color differences between the target and reference PSFs.

One way to get around both the color and normalization problems is to take images of the central source at different orientations and subtract one from the other (roll subtraction). This can be done by either requesting a roll of the telescope about the optical axis (up to 30° total) between orbits or by revisiting the object at a later date when the default orientation of the telescope is different. This technique only works when the nearby object of interest is not azimuthally extended. It is the best method for detecting point source companions or imaging strictly edge-on disks (e.g. Beta Pictoris). This method can also be used to reduce the pixel-to-pixel variations in the subtraction residuals by rotating and co-adding the images taken at different orientations (this works for extended sources if another PSF star is used). Ideally, the subtraction errors will decrease as the square root of the number of orientations.

Figure 5.11: Residual errors from the subtraction of one image of Arcturus from another taken in a different visit (filter=F435W, D=1.8" spot). The image is 29" across and has not been geometrically corrected. Logarithmic intensity scaled.


 

The large sizes of the occulting spots severely limit how close to the central source one can image. It may be useful to combine coronagraphic imaging with direct observations of the target, allowing the central columns to saturate (additional observations at other rolls would help). PSF subtraction can then be used to remove the diffracted and scattered light.

Figure 5.12: Subtraction of Arcturus from another image of itself taken during another visit using the large (D=3.0") spot and F435W filter. The image has been rebinned, smoothed, and stretched to reveal very low level residuals. The broad ring at about 13" from the star is a residual from some unknown source - perhaps it represents a zonal redistribution of light due to focus differences (breathing) between the two images. The surface brightness of this ring is 20.5 mag arcsec-2 fainter than the star. The diameter, brightness, and thickness of this ring may vary with breathing and filter. The image has not been geometrically corrected.


 

5.2.6 The Off-Spot PSF

Objects that are observed in the coronagraphic mode but that are not placed behind an occulting mask have a PSF that is defined by the Lyot stop. Because the stop effectively reduces the diameter of the telescope and introduces larger obscurations, this PSF is wider than normal, with more power in the wings and diffraction spikes (Figure 5.14). In addition, the stop and spot substrate reduce the throughput by 52.5%. In F814W, this PSF has a peak pixel containing 4.3% of the total (reduced) flux and a sharpness (including CCD charge diffusion effects) of 0.010 (compare these to 7.7% and 0.026, respectively, for the normal HRC PSF). In F435W the peak is 11% and the sharpness is 0.025 (compared to 17% and 0.051 for the normal F435W PSF). Observers need to take the reduced throughput and sharpness into account when determining detection limits for planned observations. Tiny Tim can be used to compute off-spot PSFs.

Figure 5.13: Plots of the azimuthal RMS subtraction residual levels at each radius for the large (3") spot. The flux units are counts per pixel relative to the total unocculted flux from the central source. These plots were derived from Arcturus-Arcturus subtractions represent the best results one is likely to achieve. The undistorted HRC scale assumed here is 25 mas/pixel.


 
Figure 5.14: Image of Arcturus taken in coronagraphic mode with the star placed outside of the spot. The coronagraphic field PSF has more pronounced diffraction features (rings and spikes) than the normal HRC PSF due to the effectively larger obscurations introduced by the Lyot stop. The central portion of this image is saturated. It was taken through a narrowband filter (F660N) and is not geometrically corrected


 

5.2.7 Occulting Spot Motions

Monitoring of the occulting spot positions using earth flats shows that they move over weekly, and even daily, time scales in an unpredictable manner. The cause of this motion is unknown. Their locations typically vary by ~0.3 pixel (8 mas) over the span of one week, but on occasion they have shifted by 1-5 pixels over 1-3 weeks. While inserted in the beam, however, they remain stable to better than +/-0.1 pixel, and when cycled within an orbit they return to the same position within +/-0.25 pixel.

The uncertainties in the locations lead to star-to-spot registration errors. After the acquisition exposure the star is moved to the predefined aperture position of the spot, which is measured on the ground from on-orbit flats and may be out of date. Due to the layout of ACS, it is not possible to determine the spot location automatically before a coronagraphic observation, as can be done for NICMOS. Also, unlike STIS, the star cannot be dithered around until the stellar flux within the occulter is minimized.

Star-to-spot registration errors affect coronagraphic imaging. If the star is significantly offset from the spot center (>3 pixels), then one side of the spot interior and edge will be brighter than expected and may possibly saturate much earlier than predicted. A large offset will also slightly degrade the coronagraphic suppression of the diffraction pattern. Most importantly, even slight changes in the spot position will alter the residual diffraction pattern, introducing mismatches between the target and reference PSFs that may result in large subtraction residuals. This means that an observer cannot rely on reference PSFs taken in other programs or at different times.

To reduce the impact of spot motion, observers using the ACS coronagraph are required to obtain a reference PSF in an orbit immediately before or after their science observation. A single reference PSF can be used for two science targets if all three objects can be observed in adjacent orbits and are of similar color (note that it is difficult to schedule more than five consecutive orbits). Otherwise, if multiple science targets are observed, each one will require a reference PSF. The additional reference PSF orbit(s) must be included in the Phase 1 proposal.

As of Cycle 13, STScI is using a procedure to update the coronagraphic spot positions shortly before a proposal executes to provide better registration. Currently, it takes more than three weeks to revise the official aperture location of a spot, which means that its actual position at the time of a science observation could be off by pixels. In this new procedure, the last measured offset of the spot from its defined aperture location is uploaded to HST a few orbits before a coronagraphic observation executes. By including a USE OFFSET special requirement for each coronagraphic exposure after an acquisition, the target will be shifted by the appropriate amount. The spots are measured weekly from earth flats, so this method provides more up-to-date positions than relying on the aperture location. This procedure adds approximately 40 seconds to each visit. This procedure is required for all coronagraphic observations. More details will be provided on the STScI ACS web site and in Phase II proposal instructions.

5.2.8 Planning ACS Coronagraphic Observations

Exposure Time Estimation

The estimation of exposure time for coronagraphic observations is similar to direct-mode time calculations, except that the additional background contribution from the central source's PSF has to be accounted for. Generally, most coronagraphic observations are limited by the central source's PSF wings. The APT Exposure Time Calculator includes a coronagraphic mode for estimating exposure times. We will now demonstrate how exposure times for coronagraphic observations can be determined using the web-based version of the ACS Exposure Time Calculator. The following steps are required:

Where:

In order to illustrate a calculation we shall consider the case where we are trying to determine the S/N achieved in detecting a M6V star with a V magnitude of 20.5 at a distance of 4.25 arcsec from a F0V star with a V magnitude of 6, for an exposure time of 1000 seconds with the F435W filter. Using the ACS Exposure Time Calculator and considering the case for the 3.0" occulting mask:

Observing sequence for point source companions

The best way to detect faint stellar or substellar companions is to use roll subtraction to avoid color differences between the target and reference PSFs. This also provides duplicate observations that make it easier to verify true companions from noise. It is best to roll the telescope between visits and repeat the image sequence in a new orbit. This way, you can better match the breathing cycle of the telescope than if you rolled the telescope in the middle of an orbital visibility window. You can force this to happen by selecting both orientation and time-sequencing constraints in the visit special requirements. Remember that the coronagraphic field PSF is somewhat broader than the normal HRC PSF, which may influence your assumed signal-to-noise ratio. Off-spot PSF models can be generated with the Tiny Tim PSF software. You can estimate the residual background noise level using Figure 5.13.

Suggested point-source companion observing sequence:

  1. Obtain an acquisition image
  2. Execute image sequence.
  3. Request telescope roll offset (use ORIENT 1 TO 2 FROM n special requirement in visit).
  4. Obtain another acquisition.
  5. Repeat image sequence.
  6. Repeat 3-5 as necessary.

Observing sequence for extended sources (e.g. circumstellar disks and AGN host galaxies)

When imaging extended objects, the remaining scattered light must be subtracted using a reference star image, which should match the color of your target as closely as possible. To reduce the impact of noise in the subtracted images, it helps if the reference PSF is bright enough to provide higher signal-to-noise ratios in the wings, than that of the target source. If possible, select a reference star that is nearby (<20°) and request that it be observed immediately before or after the target source. This reduces the chance that there will be large focus differences between the two visits. In order to better discriminate between subtraction artifacts and real structure, it may also help to obtain images of the target at two or more orientations of the telescope (there is no need to get reference PSF images at different rolls). You can estimate the residual background noise level using Figure 5.13.

Suggested extended source observing sequence:

  1. Obtain direct images of the science target in each filter to derive normalization factors
  2. Obtain an acquisition image of the science target.
  3. Take image sequence of science target.
  4. Request a new telescope orientation.
  5. Repeat steps 2-3.
  6. In a new visit immediately after the science observation, point at the reference star.
  7. Obtain an acquisition image of reference star.
  8. Take image sequence of reference star.
  9. Obtain direct images of the reference star in each filter to derive normalization factors.

Note that the order of the observations places direct imaging before or after coronagraphic imaging. This reduces cycling of the coronagraphic mechanism. Because the occulting spots are large, you may wish to image closer to the source using additional direct observations without the coronagraph. Multiple roll angles are necessary in this case because portions of the inner region will be affected by saturated columns and diffraction spikes. Direct observations of the reference star will be required as well to subtract both the diffracted and scattered light. Color and focus mismatches between the target and reference PSFs will be even more important in the direct imaging mode than with the coronagraph because the diffracted light is not suppressed. However, there are no mismatches caused by star-spot alignment to worry about.

5.2.9 Choice of Filters for Coronagraphic Observations

All of the HRC filters are available for coronagraphic observations. However, there are only a few that have been used in the past in this mode, and can be considered well-characterized. These are F435W, F475W, F606W, and F814W. The first three have produced good results and can be considered "safe" choices.

Filter F814W has been problematic, however. Because the PSF is larger at red wavelengths than blue, less of the light from the central star is blocked by the occulting spot. This makes the residual PSF more sensitive to spot shifts and misalignments of the star behind the spot. Mismatches between target and reference PSF star-to-spot alignments may cause significant subtraction residuals. Also, F814W images suffer from the red halo, which still affects coronagraphic observations because the central spot is filled with light that will be scattered to large radii. Color differences between the target and reference PSF stars will thus cause differences in the halo pattern, altering the local background level in unpredictable ways. If multicolor images are needed, it may be safer to choose F435W and F606W rather than F606W and F814W, for instance. Of course, for red objects it may not be feasible to choose a blue filter due to low flux levels, in which case F814W must be used.

5.3 Grism/Prism Spectroscopy

The ACS filter wheels include four dispersing elements for low resolution slitless spectrometry over the field of view of the three ACS channels. One grism (G800L) provides low resolution spectra over the 5500-11,000Å range for both the WFC and HRC; a prism (PR200L) in the HRC covers the range 1600 to beyond 3900Å; in the SBC a LiF prism covers the wavelength range 1150 to ~1800Å (PR110L) and a CaF2 prism is useful over the 1250 to ~1800Å range (PR130L). Table 5.3 summarizes the essential features of the four ACS dispersers in the five available modes. The grism provides first order spectra with dispersion almost linear with wavelength but with second order overlap beyond about 10,000Å; the prisms however have non-linear dispersion with maximum resolution at shorter wavelengths but much lower resolution at longer wavelengths. The two-pixel resolution is listed for each grism or prism at a selected wavelength in Table 5.3. The pixel scale for the prism spectra is given at the selected wavelength. The tilt of the spectra to the detector X axis (close to the spacecraft V2 axis) is also listed.


Table 5.3: Optical Parameters of ACS Dispersers
Disperser Channel Wavelength range (Å) Resolution Å/pixel Tilt1 (deg)
G800L WFC 1st order: 5500-11000 100@8000Å 39.82 -2
G800L WFC 2nd order: 5000-9500 200@8000Å 20.72 -2






G800L HRC 1st order: 5500-11000 140@8000Å 23.93 -38
G800L HRC 2nd order: 5500-9500 280@8000Å 12.03 -38
PR200L HRC 1600-3900 59@2500Å 21.3 -1






PR110L SBC 1150-1800 79@1500Å 9.5 0
PR130L SBC 1250-1800 96@1500Å 7.8 0
1Tilt with respect to the positive X-axis of the data frame.
2The dispersion varies over the field by +/- 11%; the tabulated value refers to the field center.
3The dispersion varies over the field by +/- 2%; the tabulated value refers to the field center.

5.3.1 WFC G800L

The G800L grism and the WFC provide two-pixel resolution from 69 (at 5500Å) to 138 (at 11,000Å) for first order spectra over the whole accessible field of 202x202''. Table 5.3 lists the linear dispersion, but a second order dispersion solution provides a better fit. Figure 5.15 shows the wavelength extent and sensitivity for the zeroth, first, and second order spectra when used with the WFC; Figure 5.16 shows the same plot in pixel extent. The 0 position refers to the position of the direct image and the pixel size is 0.05''. Note that there is contamination of the 1st order spectrum above 10,000Å by the second order. The total power in the zeroth order is 2.5% of that in the first order, so locating the zeroth order may not be an effective method of measuring the wavelengths of weak spectra. The default method will be to obtain a matched direct image-grism pair. There is also sensitivity of about a percent of first order in the third and fourth orders and about half a percent in the negative orders. The full extent the spectrum of a bright source (orders -2, -1, 0, 1, 2 and 3) covers is 1200 pix (60''). Figure 5.17 shows the full spectrum extent for a 60s exposure of the white dwarf GD153 (V=13.35) with the individual orders indicated. When observing bright objects, the signal in fainter orders may be mistaken for separate spectra of faint sources and in crowded fields many orders from different objects can overlap. The wavelength solution is field dependent on account of the tilt of the grism to the optical axis and the linear dispersion varies by ±11% from center to corner. This field dependence has been calibrated to allow wavelength determination to better than 0.5 pixels over the whole field.

Figure 5.15: Sensitivity versus wavelength for WFC G800L


 

5.3.2 HRC G800L

When used with the HRC, the G800L grism provides higher spatial resolution (0.028'') pixels than the WFC and also higher spectral resolution. However, the spectra are tilted at -38 degrees to the detector X axis. Figure 5.18 shows the wavelength extent and sensitivity in the zero, first and second orders, with the pixel extent shown in Figure 5.19. Figure 5.20 shows the observed spectrum of the standard star GD153. Again there is contamination of the first order spectrum by the second order beyond 9500Å. The total extent of the spectrum (orders -1 and +2) in Figure 5.20 covers about 70% of the 1024 detector pixels. In addition, a much greater number of spectra will be formed by objects situated outside the HRC direct image, or will have their spectra truncated by the chip edges, than for the WFC. The variation of the grism dispersion over the HRC field is about +/-2% from center to corner and has been calibrated.

Figure 5.16: Sensitivity versus pixel position for WFC G800L


 

5.3.3 HRC PR200L

The maximum pixel resolution of the prism is 5.3Å at 1800Å. At 3500Å, the dispersion drops to 91Å/pix and is 515Å/pix at 5000Å. The result is a bunching up of the spectrum to long wavelengths with about 8 pixels spanning 1500Å. For bright objects, this effect can lead to blooming of the HRC CCD from filled wells; the overfilled pixels bleed in the detector Y direction, and would thus affect other spectra. Figure 5.21 shows the sensitivity versus wavelength for PR200L and the wavelength extent of the pixels is indicated. The variation of the dispersion across the detector for PR200L amounts to about ±3%. The angle of the prism causes a large deviation between the position of the direct object and the region of the dispersed spectrum. The pixel numbers on Figure 5.21 indicate the size of the offset from the direct image. On account of the size of this offset, special apertures have been defined in the observation scheduling system so that the spectrum of the target centered on the direct image occurs near the center of the field in the prism image.

Figure 5.17: Full Dispersed Spectrum for White Dwarf GD153 with WFC/G800L. The numbers indicate the different grism orders.


 
Figure 5.18: Sensitivity versus wavelength for HRC G800L


 

5.3.4 SBC PR110L

The PR110L prism is sensitive to below 1200Å and includes the geo-coronal Lyman-alpha line, so it is subject to high background. The dispersion at Lyman-alpha is 2.6Å per pixel. Figure 5.22 shows the sensitivity with wavelength and the wavelength width of the pixels. The long wavelength cut-off of the CsI MAMA detector at ~1800Å occurs before the long wavelength build-up of flux; the dispersion at 1800Å is 21.6Å/pixel. However the detected counts at the long wavelength edge must be within the MAMA Bright Object Protection Limits (see Section 7.5). These limits must include the contribution of the geo-coronal Lyman-alpha flux per SBC pixel. The numbers in Figure 5.22 show the offset of the spectrum from the direct image.

5.3.5 SBC PR130L

The short wavelength cut-off of the PR130L prism at 1250Å excludes the geocoronal Lyman-alpha line, making it the disperser of choice for faint object detection in the 1250-1800Å window. The dispersion varies from 1.65Å at 1250Å to 20.2Å at 1800Å. Figure 5.23 shows the sensitivity versus wavelength and the pixel widths in angstroms. Bright Object Protection (BOP) considerations similar to the case of PR110L also apply to the use of this prism, except that the background count rate is lower (see Section 7.5).

Figure 5.19: Sensitivity versus pixel position for HRC G800L


 
Figure 5.20: Full Dispersed Spectrum of White Dwarf GD153 with HRC/G800L. The numbers indicate the different grism orders.


 

5.3.6 Observation Strategy

The default observing mode for all ACS WFC grism and HRC grism and prism modes is to obtain a direct image of the field followed by a dispersed grism/prism image. This combination will then allow the wavelength calibration of the individual target spectra by reference to the corresponding target on the direct image. The direct image will be added by default to the dispersed image by the scheduling system (AUTOIMAGE=YES). For the WFC and HRC G800L spectra, an F606W exposure will be employed and, for the HRC PR200L prism an F330W image will be applied. The companion direct image can be switched off by AUTOIMAGE=NO, allowing, for example, a separate direct image in a different filter to be specified if desired. For the SBC, the default mode will be to obtain spectra without an accompanying direct image on account of the need for Bright Object Protection (BOP), see Section 7.5. The user can separately specify a direct image to accompany the prism image, which should be scheduled immediately before or after the prism image in the same orbit. The direct image must also of course pass the BOP check.

Figure 5.21: Sensitivity versus wavelength for HRC/PR200L, numbers in figure show offset in pixels from direct image.


 

All exposures with the SBC prisms must fall within the Bright Object Protection limits. In the case of spectra, the most important determination is that the flux at the longest wavelength must not exceed 50 counts/s/pix. Table 5.4 lists for the PR110L and PR130L prisms, the observed magnitudes of stars of various spectral types whose spectra are expected to just exceed this BOP limit.


Table 5.4: BOP limits for SBC Prism spectra
Spectral Type PR110L PR130L
O5V 15.8 15.6
A1V 11.0 10.9
G2V 2.9 2.9

Figure 5.22: Sensitivity versus wavelength for SBC PR110L, numbers in figure show offset from direct image.


 

Table 5.5 lists the V detection limits for the ACS grism/prism modes for various spectral types without reddening. An exposure time of 1 hour was assumed with LOW Zodiacal background and a signal-to-noise of 5 per resolution element. For the WFC and HRC exposures, a CR-SPLIT of two was used and GAIN=1.


Table 5.5: V detection limits for the ACS Grism/Prism modes
Mode V limit for a given spectral type Wavelength of reference (Å)

O5V A1V K4V
WFC/G800L 24.4 24.6 25.5 7000
HRC/G800L 23.6 23.8 24.6 7000
HRC/PR200L 26.0 23.7 19.5 3000
SBC/PR110L 23.8 19.6 1.9 1700
SBC/PR130L 24.6 20.4 2.6 1700

Figure 5.15 through Figure 5.23 can be used to compute the detected count rate in the various orders of the grisms and prisms given the flux of the source spectra. Chapter 6 provides details of the calculations. Depending on the wavelength region, the background must also be taken into account in computing the signal-to-noise ratio. The background at each pixel consists of the sum of all the dispersed light in all the orders from the background source. For complex fields, the background consists of the dispersed spectrum of the unresolved sources; for crowded fields, overlap in the spectral direction and confusion in the direction perpendicular to the dispersion may limit the utility of the spectra.

The ACS Exposure Time Calculator supports all the available spectroscopic modes of the ACS and is available for more extensive calculations at http://apt.stsci.edu/webetc/acs/acs_spec_etc.jsp. The current version employs the on-orbit determinations of the dispersion solution and sensitivity determination where available.

For more detailed simulations of ACS spectra, an image-spectral simulator, called SLIM, is available. This tool allows synthetic target fields to be constructed and dispersed images from spectrum templates to be formed. SLIM can simulate spectra for all the ACS spectral modes. The simulator runs under Python and an executable version is available at: http://www.stecf.org/software/SLIM/SLIM10/index.html. Version 1.0 uses a Gaussian PSF but this has been found to be an adequate representation to the Tiny Tim model of the ACS PSF. A detailed description of the tool and examples of its use are given by Pirzkal et al. (ACS ISR 01-03).

5.3.7 Extraction and Calibration of Spectra

Since there is no slit in the ACS, the Point Spread Function of the target modulates the spectral resolution. In the case of extended sources it is the extension of the target in the direction of dispersion which sets the achievable resolution. Simulations show that for elliptical sources, the spectral resolution depends on the orientation of the long axis of the target to the dispersion direction and is described in more detail in Pasquali et al. (2001) http://www.stsci.edu/hst/acs/documents/isrs/isr0102.pdf. The dispersion of the grisms and prisms is well characterized, but for the wavelength zero point it is important to know the position of the target in the direct image. For the grisms, the zeroth order will generally be too weak to reliably set the wavelength zero point. Given the typical spacecraft jitter, wavelength zero points to ±0.4 pixels should be routinely achievable using the direct image taken just before or after the slitless spectrum image. Some improvement is possible in zero point assignment if the HST jitter files are employed to determine exactly the offsets between the direct and grism images.

The jitter information can be used to obtain more accurate coordinates for the center of the FOV. These in turn allow one to determine better relative offsets between the direct and the spectroscopic images.

The wavelength extent of each pixel for the WFC and HRC G800L modes in the red is small enough that fringing modulates the spectra. For the HRC, the peak-to-peak fringe amplitude is about 30% at 9500Å, similar to STIS, and is about 25% for the WFC chips. In practice, when observing point sources, and even more so when observing extended objects, the amount of detectable fringing is significantly reduced by the smoothing effect that the PSF together with the intrinsic object size have on the spectrum in the dispersion direction. Moreover, application of an optical model for the CCD fringing can allow the fringe amplitude per pixel to be reduced to below 4% peak-to-peak.

An extraction software package, aXe, is available to extract, wavelength calibrate, flat field and flux calibrate ACS grism and prism spectra. Full details can be found at: http://www.stecf.org/software/aXe/. The package is also available in STSDAS.

Figure 5.23: Sensitivity versus wavelength for SBC/PR130L


 

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