|
|
|
|
|
|
|
|
|
|
|
|
|
|
|
|
|
|
|
|
|
|
 |
 |
|
|
| Part II: ACS Data Handbook | |||||
| ||||||
6.1 Photometry
6.1.1 Photometric Systems, Units, and Zeropoints
We strongly recommend that, whenever practical, ACS photometric results be referred to a system based on its own filters. Transformation to other photometric systems is possible (see Sirianni et al., 2004, PASP submitted) but such transformations have limited precision and strongly depend on the color range, surface gravity and metallicity of the stars considered.
Three magnitude systems, based on ACS filters are commonly used, VEGAMAG, STMAG and ABMAG. They are all based on the knowledge of the throughput of the entire system (OTA + ACS CAMERA + FILTER + DETECTOR). These synthetic systems are tied to the observed standard stars. VEGAMAG is a standard-based system and therefore is subject to systematic errors tied to the calibration of the standard star. In the last decades it has become more and more common to use photometric systems such as ABMAG (Oke, J.B. 1964, ApJ 140,689) or STMAG (Koorneef, J. et. al. 1986, in Highlights of Astronomy (IAU), Vol.7, ed. J.-P. Swings, 833) which are directly related to physical units, and they are therefore much simpler and cleaner. The choice between standard-based and flux-based systems is mostly a matter of personal preference. Any new determination of the absolute efficiency of the instrument results in revised magnitudes for all three systems.
VEGAMAG
The VEGAMAG system uses Vega (
Lyr) as the standard star. The spectrum of Vega used to define this system is a composite spectrum constructed by assembling empirical and synthetic spectra together (Bohlin & Gilliland, 2004, AJ, in press). The vega magnitude of a star with flux F:
where Fvega is the calibrated spectrum of Vega in SYNPHOT. In the VEGAMAG system, by definition, Vega has zero magnitude at all wavelengths.
STMAG and ABMAG
These two similar photometric systems are flux-based. The conversion is chosen so that the magnitude in V corresponds roughly to that in the Johnson system.
In the STMAG system, the flux density is expressed per unit wavelength, while in the ABMAG system, the flux density is expressed per unit frequency. The magnitude definitions are:
where F
is expressed in erg cm-2sec-1Hz-1, and F
in erg cm-2sec-1Å-1. Another way to express these zeropoints is to say that an object with a constant F
Zeropoints
By definition the zeropoint of an instrument is the magnitude of an object that produces one count per second. Each zeropoint refers to a count rate measured in a specific aperture. Since for point source photometry it is not practical to measure counts in a very large aperture, counts are instead measured in a small aperture and an aperture correction is applied to transform the result to an "infinite" aperture. For ACS all zeropoints refer to a nominal "infinite" aperture of radius 5.5 arcsec. Independently from the different definitions, the magnitude in the passband P in any of the ACS systems is given by:
The choice of the zeropoint determines the magnitude system of ACSmag(P). There are several ways to determine the zeropoints:
- Using SYNPHOT you can renormalize a spectrum to 1 count/sec in the appropriate ACS passband and output the zeropoint in the selected magnitude system (provided that the most updated throughput tables are included in your SYNPHOT distribution). This example renormalizes a 10,000 K blackbody:
iraf> calcphot obsmode "rn(bb(10000),band(obsmode),1,counts)" form
where obsmode is the passband, for example "acs,hrc,f555w", and form is the output magnitude system: ABMAG, VEGAMAG or STMAG.
- Using the header keywords. The simplest way to calculate the most updated STMAG and ABMAG zeropoints for your data is to use the photometric keywords in the SCI extension of ACS images:
The header keywords PHOTFLAM and PHOTPLAM relate to the STMAG and ABMAG zeropoints through the formulae:
STMAG_ZPT = -2.5 Log (PHOTFLAM) - PHOTZPT = -2.5 Log (PHOTFLAM) - 21.10
ABMAG_ZPT = -2.5 Log (PHOTFLAM) - 21.10 - 5 Log (PHOTPLAM) + 18.6921
This web page will be maintained with the most current zeropoints. More information about the zeropoint determinations can be found in Sirianni et al. (2004, PASP submitted) and
ACS ISR 04-08.6.1.2 Aperture and Color Corrections
In order to reduce errors due to the residual flat fielding error and background variations and to increase the signal-to-noise ratio, both of the most popular photometric techniques, aperture photometry and PSF-fitting photometry, are usually performed by measuring the flux within a small radius around the center of the source. (For a discussion on the optimal aperture size, see Sirianni et al. 2004 PASP submitted). As a consequence, this measurement must be tied to the total count rate by applying an aperture correction.
For historic consistency, ACS zeropoints refer to the total flux in a nominal "infinite" aperture and are ready to use for surface photometry. When performing point source photometry, users should convert their count rates within any given aperture to the count rates in the infinite aperture. In order to facilitate this conversion Sirianni et al. (2004, PASP submitted) provides the correction derived from the encircled energy profiles of an intermediate aperture (0.5" radius) to an infinite aperture. This correction is filter dependent and varies between ~0.09 and ~0.12 for WFC and between ~0.09 and ~0.31 for the HRC.
Users should determine the offset between their own photometry and aperture photometry with a 0.5" radius aperture. This is usually done by measuring a few bright stars in an uncrowded region of the field of view and applying the offset to all photometric measurements. If such stars are not available, encircled energies have been tabulated (Sirianni et al. 2004, PASP submitted). However users should be reminded that accurate aperture corrections are a function of time and location on the chip and also depend on the kernel used in Drizzle. They should avoid the blind application of tabulated encircled energies especially at small radii.
The aperture correction for the near-IR filters presents further complications. All ACS CCD detectors suffer from scattered light at long wavelengths. These thinned backside illuminated devices are relatively transparent to near-IR photons. The transmitted long wavelength light illuminates and scatters in the CCD header, a soda glass substrate. It is then reflected back from the header's metallized rear surface and re-illuminates the CCDs front-side photosensitive surface (Sirianni et al. 1998, proc SPIE vol. 3555, 608, ed. S. D'Odorico). The fraction of the integrated light in the scattered light halo increases as a function of wavelength. As a consequence the PSF becomes increasingly broader at increasing wavelengths. WFC CCDs incorporate a special anti-halation aluminum layer between the frontside of the CCD and its glass substrate (Sirianni et al. 1998, proc SPIE vol.3355, 608, ed. S. D'Odorico). While this layer is effective at reducing the IR halo, it appears to give rise to a relatively strong scatter along one of the four diffraction spikes at wavelength > 9000 Å (Hartig et al. 2003, proc SPIE vol.4854, 532, ed. J.C. Blades, H.W. Siegmund), see Section 6.1.4 for more details.
The same mechanism responsible for the variation of the intensity and extension of the halo as a function of the wavelength is also responsible for the variation of the shape of the PSF as a function of color of the source. As a consequence, in the same near-IR filter the PSF for a red star is broader than the PSF of a blue star. Gilliland et al. (2002, proc SPIE vol.4854, 532, ed. J.C. Blades, H.W. Siegmund.) and Sirianni et al. (2004, PASP submitted) provide assessments of the scientific impact of these PSF artifacts in the red. The presence of the halo has the obvious effect of reducing the signal-to-noise and the limiting magnitude of the camera in the red. It also impacts the photometry in very crowded fields. The effects of the long wavelength halo should also be taken into account when performing morphological studies and performing surface photometry of extended objects (see Sirianni et al., 2004, PASP submitted, for more details).
The aperture correction for red objects should be determined using an isolated same-color star in the field of view, or using the effective wavelength versus aperture correction relation (Sirianni et al. 2004, PASP submitted). If the object's SED is available an estimate of the aperture correction is also possible with SYNPHOT. In fact, a new keyword "aper" has been implemented to call the encircled energy tables in the OBSMODE for ACS. A typical OBSMODE for an aperture of 0.5" would read like "acs,wfc,f850lp,aper$0.5". A comparison with the infinite aperture magnitude using the standard OBSMODE "acs,wfc,f850lp" would give the estimate of the aperture correction to apply. See the
SYNPHOThandbook for more details.Color Correction
In some cases it may be necessary to compare ACS photometric results with existing datasets in different photometric systems (e.g., WFPC2, SDSS or ground Johnson-Cousins). Since the ACS filters do not have exact counterparts in any other "standard" filter set, the accuracy of these transformations is limited. Moreover if the transformations are applied to objects whose spectral type (e.g., color, metallicity, surface gravity) do not match the spectral type of the calibration observation, serious systematic effects could be introduced. The transformations can be determined by using SYNPHOT, or by using the published transformation coefficients (Sirianni et al., 2004, PASP submitted). In any case users should not expect to preserve the 1-2% accuracy of ACS photometry on the transformed data.
6.1.3 Pixel Area Maps
When ACS images are flat fielded by the CALACS pipeline the resultant FLT files are flat if the original sky intensity was also flat. However, there is still very significant geometric distortion remaining in these images and the pixel area on the sky will vary around the field. Consequently, the relative photometry in the FLT images must be corrected.
One option is to drizzle the data to remove the geometric distortion and allow the sky to remain flat. Therefore both surface and point-source relative photometry are correct in the resulting DRZ files and the photometric zeropoint in the header will allow conversion from electrons/s to absolute flux units.
Unfortunately users who wish to perform photometry directly on the distorted FLT files, rather than the drizzled (DRZ) data products, will require a field-dependent correction to match their photometry with the DRZ zeropoint.
The correction to the FLT images may be made by multiplying the measured flux on the FLT image by the pixel area at the corresponding position using a pixel area map (PAM), and then dividing by the exposure time (DRZ images have units of electrons/s). The pixel area must be expressed in units of the reference scale (50mas/pixel for WFC and 25mas/pix for HRC). Once this correction has been applied the same zeropoint as applied to the drizzled data products may be used to convert to absolute flux units.
The PAM for the WFC is approximately unity at the center of the WFC2 chip, ~0.95 near the center of the WFC1 chip and ~ 1.12 near the center of the HRC.
Users may construct PAM images for themselves by downloading the script example files as well as the coefficients files which are needed from the web page:
http://www.stsci.edu/hst/acs/analysis/PAMSand executing the scripts in IRAF or PyRAF. Alternatively PAM images may be downloaded directly from the same page, but they are very large.
Figure 6.1: Variation of the WFC and HRC effective pixel area with position in detector coordinates.
PAM Concept Illustration
To illustrate the concepts of extended source and point source photometry on FLT and DRZ images we consider a simple idealized example of a 3x3 pixel section of the detector. We assume that the bias and dark corrections are zero and that the quantum efficiency is unity everywhere.
Example #1 Constant Surface Brightness Object
Let's suppose we are observing an extended object with a surface brightness of 2 e-/pixel in the undistorted case. With no geometric distortion the image is:
In reality ACS suffers from geometric distortion and as a consequence pixels are not square and the pixel area varies across the detector:
Let's suppose the pixel area map (PAM) is:
As a result in the raw data there is an apparent variation in surface brightness.
The geometrical area of each pixel is imprinted in the flat field as well as the photometric sensitivity. In this example, since we assumed that the quantum efficiency is unity everywhere, the flat field is just the equivalent of the PAM:
ACS flat fields are designed to level out a uniformly illuminated source and not to conserve total integrated counts, so after the flat field correction the FLT image has the correct surface brightness and can be used to perform surface photometry. However the image morphology is distorted.
MultiDrizzle can be run on the FLT image. The result is that each pixel is free of geometric distortion and is photometrically accurate.
Example #2 Integrated photometry of a point source
Now let's suppose we are observing a point source and that all the flux is included in the 3x3 grid. Let the counts distribution be:
The total counts are 100. Due to the geometric distortion, the PSF as seen in the raw image is distorted. The total counts are conserved, but they are redistributed on the CCD.
After the flat field correction, however, the total counts are no longer conserved:
In this example the counts now add up to 99.08, instead of 100.
In order to perform integrated photometry the pixel area variation need to be taken into account. this can be done by multiplying the FLT image by the PAM or by running MultiDrizzle.
Only by running MultiDrizzle can the geometric distortion be removed, but both approaches correctly recover the count total as 100. Users should be cautioned that this is just an idealized example. In reality the PSF of the star extends to a much bigger radius. If the user decides to work on the flat fielded image after correcting by the pixel area map, they need to calculate a new aperture correction to the total flux of the star. The aperture corrections discussed in Section 6.1.2 are only for MultiDrizzle output images. In most cases the aperture correction for distorted images will be quite different from the same star measured in the DRZ image. This is particularly true for small radius apertures.
6.1.4 PSF
PSF Field Dependence
The point spread functions in the ACS cameras are more stable over the field of view than in any other HST camera, especially when compared to WFPC2. The PSF does vary slightly in both shape and width across the field to a degree that may affect photometric measurements. This effect is described in detail in
ACS ISR 03-06. The variations in the HRC are very small and probably negligible when using apertures greater than r=1.5 pixels or using PSF fitting. However, the WFC PSF varies enough in shape and width that significant photometric errors may be introduced when using small apertures or fixed-width PSF fitting.The WFC PSF width variation is mostly due to changes in CCD charge diffusion. Charge diffusion, and thus the resulting image blur, is greater in thicker regions of the detector (the WFC CCD thickness ranges from 12.6 to 17.1 microns, see Figure 6.2). At 500 nm, the PSF FWHM varies by 25% across the field. Because charge diffusion in backside-illuminated devices like the ACS CCDs decreases with wavelength, the blurring and variations in PSF width will increase towards shorter wavelengths. At 500 nm, photometric errors of as much as 15% may result when using small (r < 1.5 pixel) apertures. At r=4 pixels, the errors are reduced to <1%. Significant errors may also be introduced when using fixed-width PSF fitting (see
ACS ISR 03-06).The shape of the PSF also changes over the WFC field due to the combined effects of aberrations like astigmatism, coma, and defocus. Astigmatism noticeably elongates the PSF cores along the edges and in the corners of the field. This may potentially alter ellipticity measurements of the bright, compact cores of small galaxies at the field edges. Coma is largely stable over most of the field and is only significant in the upper left corner, and centroid errors of ~0.15 pixels may be expected there.
Observers may use TinyTim to predict the variations in the PSF over the field of view for their particular observation. TinyTim accounts for wavelength and field-dependent charge diffusion and aberrations, and can be found at:
http://www.stsci.edu/software/tinytim/
Figure 6.2: WFC chip thickness (left) and PSF FWHM (right).
PSF Long-Wavelength Artifacts
Long wavelength (
> 700 nm) photons can pass entirely through a CCD without being detected and enter the substrate on which the detector is mounted. In the case of the ACS CCDs, the photons can be scattered to large distances (many arcseconds) within the soda glass substrate before reentering the CCD and being detected. This creates a large, diffuse halo of light surrounding an object, called the "red halo." This problem was largely solved in the WFC by applying a metal coating between the CCD and the mounting substrate that reflects photons back into the detector. Except at wavelengths longer than 900 nm (where the metal layer becomes transparent), the WFC PSF is unaffected by the red halo. The HRC CCD, however, does not have this fix and is significantly impacted by the effect.
The red halo begins to appear in the HRC at around 700 nm. It exponentially decreases in intensity with increasing radius from the source. The halo is featureless but slightly asymmetrical, with more light scattered towards the lower half of the image. By 1000 nm, it accounts for nearly 30% of the light from the source and dominates the wings of the PSF, washing out the diffraction structure. Because of its wavelength dependence, the red halo can result in different PSF light distributions within the same filter for red and blue objects. The red halo complicates photometry in red filters. In broadband filters like F814W and especially F850LP (in the WFC as well as the HRC), aperture corrections will depend on the color of the star, see Section 6.1.2. for more discussion of this. Also, in high-contrast imaging where the PSF of one star is subtracted from another (including coronagraphic imaging), color differences between the objects may lead to a significant residual over-or-under-subtracted halo.
In addition to the halo, two diffraction-spike-like streaks can be seen in both HRC and WFC data beyond 1000 nm (including F850LP). In the WFC, one streak is aligned over the left diffraction spike (right spike in the HRC) while the other is seen above the right spike (below the left spike in the HRC). These seem to be due to scattering by the electrodes on the back sides of the detectors. They are about five times brighter than the diffraction spikes and result in a fractional decrease in encircled energy. They may also produce artifacts in sharp-edged extended sources.
Figure 6.3: CCD scatter at red wavelengths. WFC images of the standard star GD71 through filters F775W (9s, left), F850LP (24s, middle), and FR1016N at 996 nm (600s, right). The CCD scatter, undetected below ~800 nm, grows rapidly with longer wavelength. In addition to the asymmetrical, horizontal feature, a weaker diagonal streak also becomes apparent near 1 micron.
HRC and SBC UV PSFs
Below 3500Å the low and mid-spatial-frequency aberrations in HST result in highly asymmetric PSF cores surrounded by a considerable halo of scattered light extending 1-2 arcseconds from the star. The asymmetries may adversely affect PSF fitting photometry if idealized PSF profiles are assumed. Also, charge scattering within the SBC MAMA detector creates a prominent halo of light extending about 1" from the star that contains roughly 20% of the light. This washes out most of the diffraction structure in the SBC PSF wings.
6.1.5 CTE
To date, all CCDs flown in the harsh radiation environment of HST suffer degradation of their charge transfer efficiency (CTE). The effect of CTE degradation is to reduce the apparent brightness of sources, requiring the application of photometric corrections to restore measured integrated counts to their "true" value.
The principle of correcting for CTE losses to achieve unbiased measurements is straightforward, but in practice it can be quite difficult to utilize. The charge loss due to imperfect CTE depends on combinations of scene characteristics such as the total counts and physical extent of a source on the detector, its background rate (global and local), the number of parallel and serial transfers, and elapsed time in the life of the detector. It is difficult to quantify and calibrate the degree to which each of these factors affects the measurements of a source over time. The general solution has been to use repeated observations of cluster fields to quantify the impact of CTE on point sources, to use internal measures of CTE degradation to track changes in between costly, external calibrations and, in the future, to do additional studies of the impact to more complicated sources (e.g., extended sources) when the CTE losses become large.
ACS has two CCDs, WFC and HRC, which clock charge and are therefore subject to CTE losses. Here we describe the external calibration of their photometric losses due to imperfect CTE. All CTE calibrations were derived from matched sets of observations of 47 Tucanae using a wide range of position, filter, and exposure time. All CTE measurements were performed by comparing the brightness of individual stars as a function of pixel transfers. The present discussion summarizes the results presented in
ACS ISR 03-09andACS ISR 04-06.Figure 6.4 gives a clear and consistent snapshot of the impact of imperfect CTE on photometry in early 2003 for WFC. The statistics are sufficient to detect non-trivial CTE losses from individual bins (~a few hundred stars) at the ~5-10 sigma confidence level for faint stars.
Figure 6.4: An example of the linear relationship determined for WFC for one filter and exposure time (F606W for 30 sec providing a sky level of <s>=3.6 electrons) at one brightness range (stars with an average of <f>~400 electrons and a full range of +/-1 mag) contained within a r=3 pixel aperture. The upper panel shows a parallel transfer relation, the lower panel illustrates the serial transfer dependence.
Figure 6.5: The dependence of parallel transfer CTE loss for WFC (over 2048 transfers) versus stellar flux contained within an r=3 pixel aperture (all images at low sky level and taken in March 2003). As shown, the relation is strong and suggests a power-law relation which was utilized in the correction formulae. The panel on the right shows the same in log-log space.
WFC Parallel and Serial CTE
As seen in Figure 6.5, the parallel WFC CTE loss appears to be a strong function of the stellar flux. This is not surprising given similar experience with WFPC2 and STIS. Charge traps may be present at varying depths of the silicon substrate of WFC. Larger charge packets (from brighter stars) may be subject to greater absolute flux losses by accessing deeper traps, but smaller charge packets (from fainter stars) appear to lose a larger fraction of their flux. As seen in Figure 6.5, the data suggests a power-law relationship between CTE loss and stellar flux which we have utilized in the following correction formula.
In Figure 6.6 we show the relationship between CTE loss and background. Increased background can mitigate the CTE losses (presumably by filling traps in advance of the arrival of the stellar charge packet during read-out), and this phenomenon has been seen for WFPC2 and STIS. For the dependence on background we also use a power-law in our correction formula.
Figure 6.6: The relationship between parallel transfer CTE loss and sky background for WFC in March 2003. As seen, the correlation is weak (and is insignificant for apertures of r=5 and 7) and in the sense of reduced loss with higher background. The measurement was made for stars with a flux of ~1000 electrons to remove the dependence on stellar flux.
The term, (Y/2048) reflects the linear relationship of the CTE loss with pixel transfer as seen in Figure 6.4. The last term involves the modified Julian date (MJD) from the header and reflects a linear degradation with time. The full range of sky values (SKY) sampled extended from ~0.1 e- (F502N for 30 sec) to 125 e- (F606W for 1100 sec). The full range of aperture flux values (FLUX) sampled ranged from ~100 e- to 300,000 e-.
Table 6.1 contains the best fit values of the parameters A, B, and C and their uncertainties for the equation given above, for 3 different aperture sizes. The parameters were determined by simultaneously fitting data obtained in March 2003, August 2003 and February 2004. The fit to the data is shown in Figure 6.7. To determine the CTE correction (YCTE) expressed in magnitudes for a particular image observers should:
- select the appropriate aperture size,
- measure the flux in the aperture and the sky value (if working with drizzled data in e-/sec one must multiply by the exposure time),
- determine the number of parallel transfers, Y (Y=y coordinate for 1<y<2048, Y=4096-y for y>2048),
- determine the MJD,
- and evaluate the correction.
The correction must be derived for the individual exposures, not for a "stack". Figure 6.7 shows the corrections derived from the formula for different observing parameters.
Table 6.1: CTE Correction Coefficients and Uncertainties for WFC Parallel Transfer.
aperture radius (pixels) A( )
B( )
C( )
In principle, first fitting the CTE trends in bins of stellar flux (and individual sky levels), followed by the step of constraining the multi-parameter fitting formula could result in information loss. However, the combination of contaminants (cosmic rays) and the large range of photon statistics (requiring different sigmas for sigma-clipping) dictated the adopted approach.
As of 2004 we have seen no evidence of photometric losses due to imperfect serial CTE on the WFC.
Figure 6.7: Parallel CTE trends. The rows and columns show the photometric losses (at y=2048) for different sky levels and dates, respectively for aperture radii of 3 pixels. The fitted line uses the global, time-dependent correction formula in "WFC Parallel and Serial CTE" with the appropriate stellar flux and sky levels.
HRC Parallel and Serial CTE
The same procedure described above for WFC was utilized to analyze HRC data and to quantify the photometric losses due to imperfect CTE. By utilizing all 4 read-out amplifiers, we obtained 2 independent measures of the CTE loss trend for all combinations of filter and exposure time and stellar flux.
Overall, the quantification of photometric losses due to imperfect CTE on HRC is of much lower accuracy than for WFC due to the small field of view of HRC (45 times less area) and the resulting paucity of stars. The adopted parallel CTE correction parameterization in magnitudes is:
The terms in the equation are the same as for the WFC equation except for the modified Julian date (MJD) term, which characterizes the time evolution of the degradation of CTE away from the calibration date (MJD=52714). Table 6.2 contains the best fit values of the parameters A, B and C and their uncertainties for the YCTE equation for HRC. At present, our calibration of CTE on HRC is based on a single epoch (March 2003).
The method to correct HRC data is the same as for WFC. Observers should:
- select the appropriate aperture size,
- measure the flux in the aperture and the sky value (if working with drizzled data in electrons/sec one must multiply by the exposure time),
- determine the number of parallel transfers, Y (Y=y coordinate for read-outs with amps C or D, Y=1024-y for read-outs with amps A or B),
- determine the MJD,
- and solve for the correction.
For most observations the default amplifier, C, is used but users can determine which amplifier was used by consulting the header keyword CCDAMP. Similar to the case for WFC, at present we see no significant photometric loss in HRC due to serial transfers.
Table 6.2: CTE Correction Coefficients and Uncertainties for HRC Parallel Transfer.
aperture (pixels) A( )
B( )
C( )
Similar to the case for WFC, at present we see no significant photometric loss in HRC due to serial transfers.
CTE Trending and Evolution
The YCTE correction formulae include a linear time-dependence term which is bounded by the condition that the photometric losses due to imperfect CTE were zero at launch (MJD=52333). While future data will help refine the time-dependence, we believe the linear time-dependence is well motivated by other, on-orbit data such as internal CTE measurements (see "Internal CTE") and measurements of cosmic ray charge deferred tails.
The Relationship between field-dependent Charge Diffusion and CTE
The CCDs in HRC and WFC are both thinned and the thinning process inevitably results in large-scale variations in the CCD thickness. Because the charge diffusion of the CCD depends on thickness (thicker regions suffer greater diffusion), the breadth of the PSF and aperture corrections can also exhibit a field-dependence (see Section 6.1.4.).
ACS ISR 03-06has characterized the spatial variation of charge diffusion as well as its impact on fixed-aperture photometry. For intermediate and large apertures (r> 4 pixels), the spatial variation of photometry is very small (less than 1%), but becomes significant for small apertures (r<3 pixels).It may be of interest to decouple the impact of imperfect CTE and charge diffusion on the field dependence of photometry but it is not necessary for observers to consider these two effects separately. The correction formulae given in Section 6.1.5 include both effects and hence corrects photometry for both effects. This is true as long as the user seeks to correct photometry to a perfect CTE and fiducial PSF width standard. An aperture correction from the fiducial aperture to a nominal "infinite" aperture is still necessary (see Section 6.1.2).
Internal CTE
Internal CTE measurements have been obtained since before launch for both the WFC and HRC. Extended Pixel Edge Response (EPER) and First Pixel Response (FPR) data is routinely collected to monitor the relative degradation of the CTE over time. These data indicate a linear degradation but are otherwise not directly useful for the scientific calibration of ACS data. These results will be published in an ISR that is presently in preparation and will be posted on the web site:
http://www.stsci.edu/hst/acs/documents/isrs/.6.1.6 Red Leak
HRC
When designing a UV filter, a high suppression of off-band transmission, particularly in the red, has to be traded with overall in-band transmission. The very high blue quantum efficiency of the HRC compared to WFPC2 makes it possible to obtain an overall red leak suppression comparable to that of the WFPC2 while using much higher transmission filters. The ratio of in-band versus total flux, determined with SYNPHOT, is given in Table 6.3 for the UV and blue HRC filters, where the cutoff point between in-band and out-of-band flux is defined as the filter's 1% transmission points. The same ratio is also listed for the equivalent filters in WFPC2. Clearly, red leaks are not a problem for F330W, F435W, and F475W and are more important for F250W and F220W. In particular, accurate UV photometry of objects with the spectrum of an M star will require correction for the red leak in F250W and will be essentially impossible in F220W. For the latter filter a red leak correction will also be necessary for K and G types.
Table 6.3: In-band Flux as a Percentage of the Total Flux.
Stellar Type WFPC2 F218W HRC F220W WFPC2 F255W HRC F250W WFPC2 F300W HRC F330W WFPC2 F439W HRC F435W WFPC2 F450W HRC F475W
SBC
The visible light rejection of the SBC is excellent, but users should be aware that stars of solar type or later will have a significant fraction of the detected flux coming from outside the nominal bandpass of the detector. Details are given below in Table 6.4.
Table 6.4: Visible-Light Rejection of the SBC F115LP Imaging Mode.
Stellar Type Percentage of all Detected Photons which have Percentage of all Detected Photons which have
6.1.7 UV Contamination
In an ongoing calibration effort, the star cluster NGC6681 has been observed since the launch of ACS to monitor the UV performance of the HRC and SBC detectors. Preliminary results for the HRC detector for the first year following launch were published in
ACS ISR 04-05. For the three filters, F220W, F250W, and F330W, 8 standard stars in the field were routinely measured, indicating a sensitivity loss of not more than ~1 to 2% per year. New data obtained in the 3 years since this analysis was performed are currently being studied and will be used to confirm this result and to characterize any dependence of loss with wavelength for the HRC.New results for the SBC detector have been obtained using repeated measurements of ~50 stars in the same cluster to provide an estimate of the UV contamination in addition to the accuracy of the existing flat fields (see ISR
ACS 2005-13). The uniformity of the detector response must be corrected prior to investigating any temporal loss in sensitivity. (See Section 5.4.2 for a discussion of the SBC L-flat corrections.)The absolute UV sensitivity just after launch appears to decline linearly for the first ~1.6 years in orbit and then to level off. Following the methodology for applying the STIS time-dependent sensitivity corrections (
Table 6.5: UV sensitivity loss for the SBC detector for the first 1.6 years in orbit. The UV sensitivity decreases linearly with time during this period and then levels off after 1.6 years, indicating no further loss.STIS ISR 2004-04), these data have been fit using line segments. Because the sensitivity losses moderate after ~1.6 years in all five SBC filters, this was chosen as the break point for the two line segments. The slope of the first line segment gives the sensitivity loss and is summarized in Table 6.5 (in percent per year). A decline in UV sensitivity of ~2-4% per year is found for the SBC detector during this time. This result agrees with the sensitivity loss derived for the STIS FUV-MAMA with the G140L filter for the first 5 years in orbit, where a 2-3% loss was observed. In the last two years of life, the rate of the STIS sensitivity loss slowed significantly compared to the previous 5 years. Similarly, after 1.6 years the ACS SBC loss rate is consistent with zero, though more follow-up data is required to confirm this result.
F115LP 1406 2.1+/--0.4 F125LP 1438 3.1+/--0.4 F140LP 1527 2.7+/--0.5 F150LP 1611 3.9+/--0.4 F165LP 1758 3.3+/--0.5
The UV sensitivity monitoring program for the HRC and SBC will be continued at a frequency of twice per year to monitor any temporal changes in response. Further development to the ACS pipeline software is required before the new sensitivity corrections can be applied within the HST on-the-fly-reprocessing (OTFR) pipeline. Until that time, the losses given in Table 6.5 must be applied manually.
Over time both the HRC and the SBC have a constant sensitivity within 1% (except the SBC filter F165LP which is within 2%). The three HRC filters F220W, F250W and F330W, do not change sensitivity by more than 1%. Presently, for the SBC the only conclusion is that any change in sensitivity is less than 2% because the precision of the current analysis is limited by the lack of an accurate L-flat field for the SBC.
|
Space Telescope Science Institute http://www.stsci.edu Voice: (410) 338-1082 help@stsci.edu |