Instrument Science Report ACS 2007-06
PHOTOMETRIC CALIBRATION OF THE ACS
CCD CAMERAS
_________________________
R. C. Bohlin
2007 June 12
ABSTRACT
The absolute flux calibration of the standard WFC and HRC filters is derived from the available
constraining observations of spectrophotometric standard stars. Values for the encircled energy (EE) of one
arcsec radius relative to an infinite aperture radius are derived for hot stars and compared to the EE for
cooler stars. The sensitivity degradation for the five year ACS lifetime is defined and used to correct the ACS
photometry before deriving revised quantum efficiency (QE) curves for the CCD detectors. F850LP
observations of red stars with WFC are made consistent with F850LP photometry for hotter stars by
decreasing the WFC QE curve longward of 9400Å. Broad band QE changes with a maximum of 2.3% for
WFC are also included in the revised QE curves for both CCD cameras. Revisions of the average filter
transmissions of up to 4% are required to bring both broad and narrow band photometry into exact
agreement with synthetic photometry from the primary white dwarf (WD) stars.
The remaining inconsistencies between WD and cooler star photometry are discussed but not resolved
because of the limitations of the present data set. Until these discrepancies are resolved, the transfer of the
fundamental WD absolute flux calibration to cooler types of F8-G0 should be regarded as uncertain by at
least 2% in the 7000-10000Å range. HST flux calibrations for spectral types cooler than G may have even
larger errors.
1. Introduction
The calibration of the ACS filter complement is based on the observations of
spectrophotometric standards and the application of synthetic photometry techniques (De
Marchi et al. 2004, DeM; Sirianni et al. 2005, S05). The transmission of the filters, the
throughput of the optics, and the quantum efficiency of the detectors are measured in the lab
and combined to form a total system throughput. The essence of the ACS absolute flux
calibration is to compare observed count rates, C, of the flux standards with the predicted
count rates, P, from the integral of the stellar flux times the total throughput for each filter.
Deviations of the observations from the predictions of the synthetic photometry are corrected
by DeM and S05 by adjusting the detector QE and the average filter transmission. In this
1
work, the above process is repeated with a more robust set of observations that include cooler
stars along with an expanded set of observations of hot white dwarf (WD) stars. This new
analysis begins with the QE results and the filter transmissions found by DeM. S05 and
references therein provide details of the instrumental design and a summary of the pipeline
data processing.
Table 1 lists the number of standard star observations, N, with ACS photometry in the two CCD
cameras, WFC and HRC. The observed count rates are measured with the standard one arcsec
radius and with the “infinite" 5.5 arcsec radius (S05). Mean corrections for the one arcsec to
infinite aperture are applied to the individual measured count rates to produce total measured
count rates. Additional corrections for changing sensitivity with time are derived and applied,
along with the sensitivity change for lower WFC CCD temperature on side2. These fully
corrected count rates, C, are compared with predictions of synthetic photometry, P, derived by
integrating the system throughput, stellar flux product over the filter bandpasses. Section 2
covers the derivation of the infinite aperture encircled energy (EE) corrections; and the changes
with time are defined in Section 3. Section 4 presents the problematic deviations of C/P from
unity for stars cooler than the WDs, where possible reasons include errors in the STIS reference
SEDs, stellar variability, EE errors, and errors in the filter red-leak estimates. CCD QE errors
and bandpass shifts of the filter transmission are investigated in section 5. Following the method
of DeM, small corrections are derived to the detector QE and to the filter transmission curves in
Section 6.
Table 1. ACS photometry of Stellar Flux Standards
Star
G191B2B(3)
GD153(3)
GD71(3)
GRW+70 5824
BD+17 4804
P330E
VB8
2M0036+18
2M0559-14
Sp.Type
DA0
DA1
DA1
DA3
sdF8
G0V
M7
L3.5
T5
V
N-WFC(1)
11.773
13.346
13.032
12.77
9.47
13.01
16.81
21.34
I=19.14
23
28
34
59
20(4)
6
21
14
6
N-HRC(2)
31
24
41
63
38
7
19
8
2
WFC observations are restricted to chip 1 and to within 8.4 arcsec of the subarray center
in order to avoid flat field uncertainties and edge effects.
(2)
HRC F220W and F250W are not included, because on-orbit L-flats are not yet available.
Star restricted to be within 6.5 arcsec of chip center.
(3)
Primary Standard, where the pure hydrogen NLTE model is used for the SED.
(4)
flt files used because of saturation. Normally, crj files along with pixel area maps are used.
crj files confirmed equivalent to drz photometry.
(1)
2
2. Total Encircled Energy
Following DeM, Bohlin & Gilliland (2004, BG), and S05, a 5.5 arcsec photometric aperture
is defined to contain all the signal from a stellar point source. Sky backgrounds are defined in
the annuli of 6-8 arcsec for WFC and 5.6-6.5 arcsec for HRC. Total stellar signals are
required to define the diffuse source, surface brightness calibration and to compare with
predicted count rates from the component throughput measurements. Aperture photometry is
computed with the IDL program apphot.pro (Landsman 2004 pers. comm.). The apphot.pro
photometry is verified against IRAF results; and photometry from the crj files plus pixel area
maps is equivalent to photometry from the drz files. For the WFC observations of the bright
BD+17 4804 (BG), the flt files must be used because of saturation of the CCD in a few
central pixels of the images. However, Gilliland (2004) demonstrated the linearity of such
saturated data to 0.1%, as long as the A-to-D converter does not saturate, i.e. gain=2 is
required.
In addition to the restrictions on the one arcsec photometry listed in the footnotes to Table 1,
EE measurements are excluded if the sub-array size is less than 1024x1024, if the star is
more than 4 arcsec from the center of the sub-array, or if the sub-array lacks overscan data.
Average EE values for each of the nine bright standard stars are displayed in Figure 1 as a
function of the filter pivot wavelength. The 3σ error-in-the-mean for the average of the four
WDs the F star, and the G star are shown as the error bars; and averages for these six stars all
agree within uncertainties. Individual EE values with 1σ formal uncertainty of >0.02 for the
six hot stars and >0.055 for the three red stars are excluded from the averages. Heavy black
lines are the polynomial fits to these averages of the six hotter stars, i.e. cubic for WFC and
fourth order for HRC. Because uncertainties in the sky level are 30x more important in a 5.5
arcsec aperture than for the standard 1 arcsec bright standard star aperture, these smooth fits
as a function of pivot wavelength average the effects of residual noise and define the
correction for the total encircled energy (EE) from radius=1 to radius=5.5 arcsec. No error
bars are shown for the narrow band filters, which are excluded from the fitted data. Green
dashed lines are the EE results from S05. Table 2 summarizes the results from the fitted
curves and compares with the results of S05.
Table 2 demonstrates excellent agreement between the new determination of the EE within
one arcsec radius and the corresponding original values from tables 3-4 of S05 for WD stars.
The total count rates are defined as the response in a 5.5 arcsec radius aperture, even though
this photometry is exquisitely sensitive to the measured background sky level.
3
Fig. 1. Encircled energy, i.e. fraction of signal in a one arcsec aperture relative to an "infinite" aperture of 5.5 arcsec
radius. WFC is inset into the HRC plot. Symbols are open black square: G191B2B, filled black square: GD153, open
black triangle: GD71, blue diamond: BD+17 4708, black X: GRW+70 5824, green circle: P330E, upside down red
triangle: VB8, red square: 2M0036+18, red circle: 2M0559-14. The solid black lines are polynomial fits to the average
of the EE from the WDs, the F star (BD+17 4708), and the G star (P330E). Error bars are ±3σ error in the mean of
these averages. The green dashed lines are the results of S05.
4
Table 2. One Arcsec Encircled Energy Fractions for Hot Stars
Filter
F220W
F250W
F330W
F344N
F435W
F475W
F502N
F550M
F555W
F606W
F625W
F658N
F660N
F775W
F814W
F892N
F850LP
WFC
New
S05
...
...
...
...
0.942
0.942
0.943
0.944
0.944
0.945
0.946
0.947
0.947
0.946
0.945
0.938
0.936
...
...
...
...
0.940
0.952
0.950
0.948
0.948
0.952
0.952
0.953
0.953
0.955
0.955
0.947
0.946
HRC
New
S05
0 952
0.951
0.950
0.950
0.950
0.951
0.951
0.953
0.952
0.953
0.952
0.951
0.950
0.929
0.912
0.852
0.832
0.952
0.951
0.942
0.942
0.948
0.947
0.949
0.953
0.953
0.952
0.950
0.943
0.943
0.929
0.903
0.831
0.823
For WFC, the new results are systematically lower than S05 by <1%, despite the mostly
different data sets used to derive the EE. S05 used only GD71 and GRW+70 5824, while the
G star P330E and several more recent observations of the three primary WDs G191B2B,
GD71, and GD153 are now included. Most of the early data used by S05 are at the WFC1
reference position and are now excluded because of the lack of any overscan on the
subarrarys and consequent uncertainty in the sky level. The new WFC EEs are lower than the
S05 values, probably because S05 included the chip 2 data for GD71, which is systematically
higher by 1-4%. These chip 2 data are from small 512x512 sub-arrays with no overscan,
where the 8 arcsec outer background radius often falls outside the sub-array.
For HRC, the original data set is supplemented by many more observations of the WDs,
P330E, and the F star BD+17 4804. The only HRC difference with S05 of more than 1% is
for F892N with about 2% higher EE. The "New" values in Table 2 are from the Figure 1 fits
evaluated at the pivot wavelength of each filter. The mean and error-in-the-mean for the nine
HRC observations with F892N are 0.846+/-0.003, ie within 2σ of the 0.852 in Table 2.
Typical formal 1σ uncertainties on the fitted values in Table 2 are 0.4% for WFC and 0.6%
for HRC.
5
Red stars of spectral type M and later have a "red halo" and show significantly lower EE,
because long wavelength photons scatter more in the CCD substrates, especially for HRC,
which lacks the special anti-scattering layer incorporated into the WFC CCDs. Figure 1
shows these red star EE values as red data points. S05 provide their tables 6-7, which define
the EE parameterized by the effective wavelengths, which are not constant for a filter but
increase as the peak of the stellar flux distributions move to longer wavelengths. The S05
values for the effective wavelengths are wrong due to an error in the Synphot software used
by S05; however, those incorrect values must still be used to interpolate in the S05 tables 6-7.
Table 3 includes these S05 EE values to compare with the Table 2 values for hotter stars and
with the new values that are derived directly from the measurements.
Considering scatter among the EE measures from star to star, many of the red points are
consistent with the fits for the hotter stars copied from Table 2 to column two of Table 3. For
WFC, the only discrepancy with S05 is for 2M00559-14 in the F814W filter. There are two
direct measures of this EE, 0.886 and 0.907, both of which are more than 3σ below the S05
value of 0.94. The directly measured values for WFC F814W for 2M00559-14 and for all
three stars in F850LP in Table 3 are adopted, whereas the Table 2 values are adopted for all
other WFC cases.
For HRC, again 2M00559-14 in the F814W filter is the biggest discrepancy between the new
measurements and S05. In this case, the single measurement is fairly long at 300s and has a
formal 3σ uncertainty of 0.14. For 2M0036+18 in F892N, the single exposure time is 360s
and the 3σ uncertainty is 0.25, so the S05 value of 0.84 is preferred. Thus, all of the
smoothly parameterized EE of S05 in Table 3 are adopted for HRC, while the Table 2 HRC
value are adopted for the red stars for filters shortward of F775W. EE values in Table 3 that
are recommended in lieu of the Table 3 values are flagged.
Table 3. Encircled Energy for One Arcsec for Cool Stars
Filter Hot-Star
WFC
F814W
F892N
F850LP
0.945
0.938
0.936
VB8(M7)
New
S05
0.943 0.95
...
0.94
(1)
0.922 0.93
2M0036(L3.5)
New
S05
2M0559(T6.5)
New
S05
0.960 0.95
0.942 0.94
0.920(1)0.92
0.896(1)0.94
...
0.94
(1)
0.879
0.88
HRC
F775W
0.929
0.920 0.91(1) 0.934 0.91(1) ...
0.90(1)
F814W
0.912
0.870 0.86(1) 0.860 0.85(1) 0.608 0.82(1)
F892N
0.852
0.855 0.84(1) 0.977 0.84(1) ...
0.84(1)
F850LP 0.832
0.807 0.78(1) 0.799 0.76(1) 0.701 0.66(1)
(1)
Recommended values. Use Table 2 values for all other cases.
6
Uncertainties for the EE of red stars in HRC are larger than for the ~1% estimated for WFC.
Judging from the differences between the "New" value and the S05 values, systematic
uncertainties in HRC may be ~4% for VB8 and 2M0036+18, while the 2M0559-14
uncertainties may be 10-20%. The S05 figure 8 shows a value for VB8 in HRC F850LP
closer to the "New" value than to their interpolated value of 0.78.
The uncertainties on the Table 2 EE do not enter in the final uncertainties of ACS CCD
absolute zero points, except for the reddest stars in the long wavelength filters where the EE
is a function of the SED. Otherwise, any error in the EE is compensated by the QE and filter
transmission adjustments discussed in sections 5-6. In other words, the high precision one
arcsec aperture photometry is the primary reference for the ACS absolute calibration.
3. Sensitivity Change with Time
3.1 Lower WFC Operating Temperature after 2006 July 4
Following the switch to the side 2 electronics and the drop in the CCD temperature from -77
to -80C on 2006 July 4, the decrease in sensitivity for each of the WFC filters in mag is listed
in table 5 of Mack (2007). Additional uncertainty in making this correction to the few
observations obtained after 2006 July 4 is <~0.002 mag.
3.2 Possible CTE Losses
To date, no comprehensive study of CCD charge transfer efficiency (CTE) losses has been
done for the ACS CCDs. The estimates of CTE loss by Riess and Mack (2004) for WFC are
inadequate for this study, because: a) there is no analysis of CTE loss for the last three years
of the five year (2002.2-2007.07) WFC and HRC lifetimes when CTE losses were the
greatest. b) The maximum photometric aperture studied is seven pixels, while a radius for
WFC of 20 pixels is used for bright standard stars. c) The algorithm of Riess & Mack in their
equation 2 is proportional to the sky value to a negative power, which is undefined for the
nominal sky levels near and below zero that are common for the few second exposure time of
the bright stars in Table 1.
Nevertheless, the Riess & Mack algorithm should provide an upper limit to the CTE loss,
because these losses decrease for the larger pixel radius. For example for the faintest hot star
GD153, the minimum WFC sky is ~2 electrons, the typical flux in electrons for the cr-split=2
observations is ~250,000 electrons, and the latest times are <5 years after ACS activation at
2002.2, when the distance of the star from the readout amp is ~512px. Plugging these values
into the Riess & Mack equation 2 for the parameter values in their table 2 for a seven pixel
radius yields a maximum CTE loss of 0.2%. For HRC, a similar formulation with the sky to a
negative power is even less appropriate, because of more frequent negative sky values.
7
The three red stars in Table 1 are faint at some wavelengths; but longer integration times
keep the signal high and raise the background, so that the largest predicted CTE loss for the
WFC data is 0.4% for 2M0559-14 in F625W. In HRC, the worst case for 2M0036+18 is
F606W at 2005.5 with a sky of 1.9 and a flux of 36600 electrons for which the WFC formula
predicts a 0.5% CTE loss in a seven pixel radius. Except for 2M0559-14 and, perhaps, the
F606W and F625W filters of 2M0036+18 with HRC, CTE corrections in the 20 (WFC) and
40 (HRC) pixel radius apertures can be neglected in the following analyses.
3.3 Throughput Degradation
In order to determine sensitivity degradation over the five year lifetime of the CCD modes on
ACS, plots like the example for F555W in Figure 2 are constructed for each filter. The side 2
temperature correction has been made to the two WFC points after 2006.5. All stars are
Fig. 2. Ratio of observed to predicted count rates C/P for F555W as a function of time. Symbols are: open
square-G191B2B, filled square-GD153, open triangle-GD71, diamond-BD+17 4708, X-GRW+70 5824,
circle-P330E, upside down triangle-VB8. Black symbols are for WFC, while blue indicates HRC. Only the
large symbols for the WD stars are used to define the linear fits. Vertical offsets of the fits from unity appear
because the small QE errors discussed in sections 5-6 are not yet included.
8
shown, but only the three primary WD stars (large symbols) that cover multi-year time spans
are included in the least-square linear fits: black for WFC and blue for HRC. Systematic
offsets of the C/P values of the observed/synthetic predictions from unity are due to small
errors in the QE+filter transmission values discussed in the next sections.
Fig. 3. Slopes in percent loss of sensitivity per year from fits for all filters, as in the example for F555W in
Figure 2. Black is for WFC, while blue represents the HRC. A fourth order fit to the large symbols for the
broadband filters is shown as a function of the filter pivot wavelength. The small symbols for the narrow
band filters are not used in the fit. Error bars are the 1σ formal uncertainty in the slope of the linear fits.
Most of the error bars cross the fit, and there is no systematic difference between the WFC and HRC data
points. The dotted line is the corresponding average loss of sensitivity for the three CCD grating modes of
STIS from (Stys, Bohlin, & Goudfrooij 2004). The weight of the WFC F850LP data point is arbitrarily
increased in order to avoid fit values greater than zero.
The set of slopes in percent loss per year appear in Figure 3, where black squares are WFC
and blue circles are HRC. A fourth order polynomial (heavy line) fits the large circles and
squares. Error bars are from the uncertainty in the slope coefficient. Small symbols represent
the five narrow band filters, which are not included in the fits. Narrow band photometry may
suffer additional errors due to poor coverage over time, to spectral lines that are not perfectly
9
matched to the filter bandpass, or to poor flat fields in the case of F892N that has fringing
problems. The fit coefficients are -0.04183, 2.934e-05, -7.799e-09, 8.786e-13, and -3.522e17. The fit crosses most of the 1σ error bars and is within 1.1σ of the remainder.
Also shown in Figure 3 as a heavy dotted line is the corresponding %/yr loss for STIS,
averaged over its first five years on-orbit. The discontinuities between adjacent STIS grating
modes are probably due in part to grating blaze shifts caused by shrinkage of the grating
substrates. However, the overall similarity between STIS and ACS suggests a common
source for the degradation of throughput. Possibilities for the dominant source of sensitivity
loss include the primary mirror, which is directly exposed to environmental radiation.
Another single component possibility is inside the detector housing, if the contamination
sources were sealed inside the detector modules. Usually, HST sensitivity losses are
attributed to polymerization of hydrocarbons on the optics within the scientific instruments
themselves, which would imply that losses should be proportional to the number of optical
surfaces within each instrument. If internal instrument contamination scenarios are correct
and if all new instruments have similar initial out-gassing rates, then comparing the first five
years of ACS to the first five years of STIS is the appropriate comparison.
4. Residual Errors in the ACS Absolute Flux Calibration
The one arcsec photometry is corrected for the EE and for the changes with time to define
final observed count rates, C. These C values are compared with the predictions P from
synthetic photometry, as derived from the stellar absolute flux distributions and the
instrumental component throughput measurements. This procedure was used by DeM to
provide the original WD based calibration for the ACS filters. Since the original DeM
calibration, changes in the flat fields, the encircled energy corrections, and a mostly new set
of observations of the prime WD standards require small adjustments to the DeM calibration.
These adjustments are described in detail in section 6, while some limitations on the
precision of the measurements are presented next.
4.1 Residuals for Stars Cooler than the WDs
After making the above corrections for the side 2 temperatures, the EE, and changes with
time, the average ACS photometry C for the three prime WDs show small differences with
the synthetic photometry P from the pure hydrogen model SEDs. However, residuals in C/P
for the cooler stars differ significantly from those for the WDs. For the broadband filters,
Table 4 shows these cool star residuals relative to the WDs, i.e. the table entries are the
photometric errors after adjusting all photometry for the small average errors from the three
prime WDs. Tabulated values are 100(C/P-1),so that negative values represent cases where
the predicted synthetic count rate P is larger than the observed C.
10
The precision of the ACS photometry is generally limited by the typical 1σ repeatability of
~0.3%, because the Poisson statistical error is often smaller at ~0.1% for the typical total
counts of about a million. A few observations where the formal uncertainties are larger than
the deviation from unity are omitted from the Table. For the stars with multiple observations,
the typical 3σ errors-in-the-mean for the Table 4 differences are 0.5%, based on 0.3-0.4% for
each of the WD and cooler star uncertainties. As relevant, larger uncertainties for a few cases
are discussed below and flagged with an asterisk for unreliable and with the uncertainty in
parentheses for useful values where the uncertainty is larger than 0.5%.
Table 4. Percent Residuals in C/P for Cooler Stars
BD+17 4708
P330E
F330W
WFC
...
...
HRC
-1.9
...
F435W
WFC
...
...
HRC
0.2
...
F475W
WFC
-0.3
...
HRC
-0.6
-0.3
F555W
WFC
-0.5
-0.5
HRC
-0.4
-0.2
F606W
WFC
...
-0.1
HRC
-0.3
-0.2
F625W
WFC
-1.7
-0.6
HRC
-1.3
-1.0
F775W
WFC
-1.7
-1.5
HRC
-1.4
-0.9
F814W
WFC
-1.6
-1.4
HRC
-2.5
-1.7
F850LP
WFC
-2.1
-1.2
HRC
-2.2
-1.4
* Large uncertainty. See below.
VB8
2M0036+18
2M0559-14
...
...
...
...
...
...
11.3*
8.9*
...
...
...
...
0.2*
...
...
...
...
...
-1.6*
-1.3*
...
...
...
...
-0.0*
-1.1*
...
...
...
...
0.9*
-2.2*
-3.6(2.0)
-5.2(2.0)
-80.0*
...
-2.9(1)
1.0*
-4.3(>1)
0.4*
-25.1*
...
-0.9(1)
-0.4*
-1.0(>1)
-0.7*
2.7*
-2.6*
-1.5(1)
1.6*
-2.2(>1)
2.9*
-2.9(1.3)
17.4*
11
4.2 STIS Background Subtraction Errors
The STIS background is measured at the large distance of 300 pixels from the spectrum in
order to avoid any contribution from the broad wings of the PSF. However, this choice makes
any fluctuation of the background perpendicular to the dispersion an important error source
in P for the lowest signal levels of the faintest stars, 2M0036+18 and 2M0559-14. The STIS
background is rather noisy because of residuals from cosmic ray tails that are not removed by
the cosmic ray rejection software. An estimate of the uncertainty is 0.005e-/s in the background level of the CCD spectra. Because this uncertainty is less important at the higher
signals at the longer wavelengths, the synthetic photometry P values are more reliable for the
longer wavelength ACS filters. Table 5 summarizes the effect of the STIS background
uncertainty for the two faintest red stars in all the observed ACS filters. The WFC and HRC
have the same uncertainty to <0.1%, because the same filter throughput is used and the QE
curves are similar. VB8 is bright enough so that the only added uncertainty bigger than 0.5%
is +/-1.4% in F435W. The large STIS uncertainties for 2M0559-14 require three filters to be
flagged as useless in Table 4, while the F850LP measurement has a large uncertainty of
1.3%.
Table 5. Errors (%) Due to STIS Background
F625W
F775W
F814W
F850LP
2M0036+18 2M0559-14
2.0
42.7
0.4
10.9
0.3
2.8
0.2
1.3
4.3 Stellar Variability in the Coolest Three Stars
At the short wavelengths below the peak of the SED, VB8 may be variable in analogy with
the solar UV variability. Figure 4 shows the repeated ACS observations as a function of time.
The first HRC observation is normalized to the WFC locus for each filter to avoid questions
of differential HRC/WFC corrections. For the five HRC filters with observations at 2003.723
and 2005.331, the worst case agreement of HRC at 2005.331 with the set of WFC
observations at 2005.359 is 0.4% for F555W and suggests that VB8 changes slowly enough
that the interpolated WFC values at 2003.723 are valid. However, this agreement in relative
brightness at these two times may be a fortuitous sampling of more rapid variability. The
largest Poisson uncertainty among all WFC and HRC observations in the seven filters of
Figure 4 is 0.2%. For the three primary WDs in F435W plus the seven filters shown in Figure
4, there are the eight filters times two cameras times three stars for a total of 48 cases, where
the maximum 3σ rms is 0.7% in F850LP for seven HRC observations that also represent the
only case where the total range of scatter exceeds 1%. Thus for VB8, the total range for
F475W (3.2%), F555W (1.7%), and F606W (2.9%) must be attributed to stellar variability.
12
F435W at a shorter wavelength is likely to vary even more than the longer wavelength filters.
These four C/P cases are flagged as unreliable in Table 4.
A small mystery remains of why the HRC renormalization value for F625W in Figure 4 is as
large as 1.031, while a renormalization of only 1.003 is required for HRC F606W. For
F625W, the similarity of the drop of ~1% for both of the last WFC and HRC data points is
also probably a symptom of variability, so that the especially anomalous 3.1% difference
between the WFC and HRC values in Table 4 for F625W is flagged as not reliable. For the
three longer wavelength filters F775W, F814W, and F850LP, the largest range among the
three measures for each camera is 0.5% for HRC F775W.
Fig. 4. ACS photometry C divided by STIS synthetic photometry P vs. time for VB8. Solid lines and filled
circles are WFC data, while dotted lines and open circles are the HRC data normalized to the WFC data as
interpolated to the time of the first HRC observation. There are no HRC data for F475W and only two HRC
observations with F606W and F625W.
For 2M0036+18, there are indications of variability in the sparse ACS data set. For example,
after 2002 Oct, where there are WFC and HRC observations within two days of each other,
the WFC F850LP shows a drop of 1% in 2006, while HRC shows a rise of 1% in 2005. Thus,
the 2M0036+18 values in Table 4 must have at least a 1% uncertainty.
13
ACS observations of 2M0559-14 were made at only one epoch for HRC and two epochs for
WFC. The repeated WFC F850LP photometry changes by 1%, i.e. four times the Poisson
uncertainty, so that the Table 4 values for 2M0559-14 must also have a minimum uncertainty
of 1% due to possible stellar variability.
4.4 Uncertainty in EE Corrections
In Table 4, the differences between WFC and HRC for the same star and the same filter are
the most problematic. In particular, the 3-5% differences between the WFC and HRC
residuals for VB8 and 2M0036+18 for F775W and F850LP are more than can be attributed to
the estimated 1% stellar variability discussed above. Errors in the STIS spectrum do not
affect this difference in C/P, because both cameras are equally affected by any error in
computing P from the STIS SED. Uncertainties in the EE energy corrections are estimated in
section 2 for the three red stars in Table 3 as 1% for WFC and large enough for HRC to flag
F775W, F814W, and F850LP as not useful in Table 4. These uncertainties compromise the
Fig. 5. Fractional contribution to the F625W signal in each of 20 equal bins spanning the full range of the
filter transmission measurements for WFC (solid lines) and HRC (dotted lines). Wavelength bins that
contribute less than 0.0001 (0.01%) are off scale low and do not appear on the plot. Colors and symbols are:
black squares-WDs, blue circles-BD+17 4804 (coincident with P330E), green circles - P330E, red inverted
triangles - VB8, red squares - 2M0036+18, red circles - 2M0559-14.
14
use of the three red stars for revising the ACS filter throughput. A much larger set of EE
measurements, especially for HRC, would be required to make their photometry more
accurate for the long wavelength filters that suffer from the red halo effect.
4.5 Out-of-Band Filter Transmission Errors
As seen Figure 5 around 10,000A, out-of-band transmission contributes slightly to the
measured 2M0559-14 signal. This contribution is ~0.01% for each of two points, so that even
if 10 of the 20 sample points contributed near this level, the total is predicted at <0.1% for the
reddest star. To make even a 1% contribution to the signal, the measured out-of-band
rejection would need to be 10x higher than measured before launch. Since the F625W
rejection is typical of the other broadband filters, out-of-band transmission errors are not a
likely source of the photometry discrepancies.
4.6 Non-linearity
Plots of C/P vs. count rate show no obvious trend that would suggest a count rate dependent
non-linearity. Gilliland (2004) has already demonstrated the excellent linearity of the ACS
CCDs as a function of signal level.
4.7 Scattered Light from the STIS Gratings
Ruled gratings can have a general scattered light background (eg. Bohlin 1975), which
should be characterized in the lab using lasers (Zong et al. 2006). On orbit, the amount of
any such contamination is difficult to measure in the absence of strong, damped interstellar
absorption lines that have zero starlight in the center of the line profiles. However, the result
of subtracting a fraction of the broad band stellar signal from the STIS G750L count rates
would be to increase the F and G star fluxes and their predicted P values with respect to the
WDs. Thus, any scattered light correction would only make the negative residuals in Table 4
for BD+17 4708 and P330E worse. The ACS photometry is consistent with the current
assumption of a negligible contribution from STIS scattered light.
5. Component Throughput Updates Required for the Cool Stars
What causes the systematic residuals in Table 4 for the cooler stars with respect to the
calibration defined by WDs? Possibilities include QE errors for the CCD detectors and filter
bandpass errors. The most difficult errors to explain are those where the WFC differs
significantly from the HRC residuals for the same star and the same filter. Fortunately, these
cases of 3-5% differences that include VB8-F625W and both VB8 and 2M0036+18 for
F775W and F850LP are flagged as unreliable due to stellar variability or uncertain EE
corrections for HRC. Otherwise, variability of the transmission with position on the filters,
which are the same component for WFC and HRC, might be required.
15
5.1 QE Adjustments
In general, adjustment of the individual detector QE cannot remove even the red star
residuals without introducing excessive structure into the measured QE curves. However, for
the unique case of WFC F850LP, each red star has a negative residual that is roughly twice
the uncertainty. The agreement within this set of three stars justifies an adjustment to the long
wavelength QE curve for WFC. Confining the QE adjustment to wavelengths longer than
9400Å minimizes the effect of QE changes on other WFC filters. The new QE curve is
compared to the current version in Figure 6 and the new residuals for the two longest
wavelength broadband filters are compared to the original residuals in Table 6. F814W
shows little change, while the F850LP residuals for VB8 and 2M0036+18 are ~0. The
2M0559-14 residual for F850LP is consistent with zero given the 1.3% uncertainty for STIS
from Table 5. The long wavelength cutoff of the WFC QE curve is set to 10924A,
corresponding the expectation for a temperature of -77C (Sirianni, private comm.)
Fig. 6. WFC QE. Dash-original pre-launch, solid black-DeM, red-new. The differences beyond 9400Å are
explained in section 5.1, and the shorter wavelength adjustments are described in section 6.
16
Table 6. Residuals in C/P for Long Wave QE Change
Star
Original-WFC
(%)
New-WFC
(%)
F814W
BD+17 4804
P330E
VB8
2M0036+18
-1.6
-1.4
-0.9
-1.0
-1.6
-1.4
-0.8
-0.8
F850LP
BD+17 4804
P330E
VB8
2M0036+18
2M0559-14
-2.1
-1.2
-1.5
-2.2
-2.9
-1.8
-0.9
-0.1
0.0
1.3
The remaining substantial average residual of -1.4% in F850LP for the F and G stars BD+17
4804 and P330E cannot be removed by adjusting the QE at wavelengths below the 9400Å
limit imposed above, because of the tiny difference in the WD vs. F-G contributions to the
integral filter response over the remaining 8000-9400Å range of the in band response. Any
attempt to correct for the F-G star errors in F850LP would require the full 8000-10500A in
band range to have enough leverage to make a half-way smooth QE correction. Even then,
the required QE adjustment would be much more than proposed to remedy the red stars error
and would, therefore, make a very large positive red star residual. Furthermore, correcting the
QE below 9400Å affects the F814W residuals and would require QE adjustments at
wavelengths below 8000Å. Then, F775W is affected and so on, until the required changes
become a hopelessly complicated morass. At best, the red star residuals in shorter wavelength
filters can be removed but without much more improvement for the F-G stars than for
F850LP in Table 6. Another negative is that such broad band trial QE modification over the
whole wavelength range longward of 6500Å makes a rather suspiciously lumpy final QE
curve.
5.2 Filter Shifts
The remaining significant negative residuals in Table 4 include many filters for the F and G
star, F625W for 2M0036+18 in both cameras, and WFC F775W for VB8 and 2M0036+18.
The ACS WFC photometry for VB8 in F775W is robust with agreement to 0.001 among all
three corrected values from 2002, 2004, and 2005.
Figures 7-8 illustrate the systematic residuals for the F star BD+17 4804 and the G star
P330E. Independent WFC and HRC photometry agree to ~0.4% for the four pairs of multiple
observations with large symbols in Figure 7 and to <0.6% for the six pairs of the single
17
Fig. 7. C/P for BD+17 4804 for each WFC (black) and HRC (blue) filter showing a systematic slope with
pivot wavelength. Small symbols are narrow band and single observations. Large symbols are multiple
observations with error bars of 3σ error in the mean. The heavy black line is a linear fit to the large symbols,
excluding F330W.
Fig. 8. For P330E as in Figure 7, except that all observations are single and are all used to define the fit.
18
Fig. 9. Relative response in filter F775W as the filter profile is shifted in wavelength by the amount shown
on the x-axis. Black line and symbols are for the WDs, while blue and green for BD+17 4708 and P330E are
nearly indistinguishable. Red inverted triangles, red squares, and red circles represent VB8, 2M0036+18, and
2M0559-14, respectively. The relative response is unity for all filters at zero shift of their currently accepted
bandpass.
Fig. 10. Absolute flux SEDs of the stars used to construct Figure 9 with the same color coding as used for
Figure 9. The dotted black line is GRW+70 5824, while the prime WD GD153 is the solid black line. The
red curves in order of bright to faint are VB8, 2M0036+18, and 2M0559-14.
19
observations of P330E in Figure 8. Lacking the G and cooler star data, BG suggested that the
residuals for BD+17 4708 with respect to the WDs could be explained by a shift in
wavelength of the filter transmissions. For example for F775W, Figure 9 shows how the
photometry changes with a shift of the transmission profile toward shorter wavelengths. The
WDs (black curves) have an increasing predicted response P for negative shifts, the F and G
stars (blue and green) are more nearly flat, and the cooler stars (red) show decreasing P
values. Plots like Figure 9 are readily understood in terms of the stellar flux SEDs shown in
Figure 10. For all the relevant broadband filters longward of 5000Å, the photometric
variation vs. shift for the F and G star are always similar to each other just as in Figure 9 for
F775W, while the predicted photometry P for the coolest stars always drops with
increasingly negative shifts. Thus, for each star with a measured error with respect to the
WDs, the required shift can be read from a Figure 9 style plot.
Following the BG procedure for BD+17 4804 (sdF8), but also including the G star P330E,
the average residuals weighted by the number of valid WFC or HRC measurements are
computed for the five broadband filters with residuals greater than the 3σ uncertainty of
~0.5%. Table 7 includes these averages (Avg-Res), the shift in Ångstroms required to reduce
that residual to zero, and the original shifts computed from a smaller data set by BG. For
WFC F850LP, the original QE curve is used, rather than the revisions of section 5.1, i.e. a
shift is an alternative to a QE revision. The last three lines of Table 7 are the percent residuals
for the red stars, where the only valid HRC measure included is for 2M0036+18 F625W. The
slash in these three lines separates the average residual from Table 4 from the residual after
applying the shift in line two.
The increase in the shift required over that proposed by BG in 2004 is largely due to the
improved CTE correction for STIS (Goudfrooij & Bohlin 2006). The STIS CTE correction is
difficult to define precisely at the longer wavelengths; and the changes in the calculated shifts
might be interpreted as the limit to the ability of the STIS spectra to provide precise synthetic
photometry P values.
One clear result in Table 7 is the increase of the residuals for the three longer wavelength
filters from small negative values to large positive residuals, which are well outside the
uncertainties of 1-2% for these measurements. Thus, a simple bulk shift of the filter
bandpasses can remedy the systematically negative residuals for the F and G stars but only at
the expensive of making the red star residuals unreasonably large for three filters. Of course,
introducing one more correction parameter by making different shifts at the short and long
wavelength filter cutoff wavelengths can provide a means of adjusting the F-G residuals
independently of the red star residuals. Because the red stars have SEDs increasing toward
longer wavelengths, their residuals depend more on the long wavelength cutoff. Separate
shifts for the short wavelength cutoff of somewhat more than the values in line 2 of Table 7
and somewhat less than those values at the long wavelength cutoff can bring all residual to
zero within uncertainties. Encouragingly, the errors for the WFC F775W and F814W filters
20
are eliminated with about the same wavelength shifts for both VB8 and 2M0036+18, while
all three stars call for about a -25Å shift of the long wave cutoff for F850LP.
Table 7. Wavelength Shifts Required by the F and G Stars
Filter
F330W
Avg-Res(%)
Shift(Å)
BG Shift(Å)
VB8(%)
2M0036+18(%)
2M0559-14(%)
-1.9
-14
-12
...
...
...
F625W
F775W
F814W
-1.1
-27
-24
...
-4.4/+1.5
...
-1.5
-57
-44
-2.9/+4.2
-4.3/+4.6
...
-1.8
-71
-55
-0.9/+7.7
-1.0/+10
...
F850LP
-2.0
-94
-60
-1.5/+5.4
-2.2/+7.0
-2.9/+12
However, to justify the implementation of such drastic two parameter modifications to the
filter throughputs requires a more robust data base of measurements than is currently
available. More red star observations would be required to fully rule out larger than expected
errors from stellar variability, EE corrections, and even possible filter transmission nonuniformities with detector coordinate. Significant errors in the P values from uncertainties in
the STIS CTE correction may also be present. Currently, the shifts in Table 7 should be
viewed as covering the range of uncertainty in our knowledge of ACS filter transmission
characteristics.
By omitting any adjustments of the wavelength dependence of the filter bandpasses, ordinary
F and G standard stars retain the surprisingly large color residuals of Tables 4 & 6 and as
illustrated in Figures 7-8. Such large systematics compromise efforts to precisely relate ACS
observations to physical properties of stars through comparison with synthetic photometry.
For example, the color F555W - F814W shows a ~0.01 mag systematic residual error for
these common F-G stars relative to the WDs, which are the basis of HST flux calibrations.
6. Traditional Adjustments for a Pure WD Flux Calibration
Adopting the new EE values of section 2, the correction for lower WFC operating
temperature after 2006 July 4 from section 3.1, the correction for loss of sensitivity with time
21
Fig.11. C/P count rate ratios C/P vs. pivot wavelength for the WD stars observed with WFC. Squares are
averages of the three prime WDs G191B2B, GD153, and GD71 with their 3σ error in the mean, while the
red Xs represents the slightly cooler GRW+70 5824. Following DeM, a smooth quadratic fit (heavy line) to
the broadband filters (filled squares) is made, ignoring the more deviant medium and narrow band filters
(open squares).
Fig.12. As in Figure 11 for HRC.
22
in section 3.3, and the WFC QE change detailed in section 5.1, the revised ACS C values are
compared with the P values for the WD stars in Figures 11-12 for WFC and HRC,
respectively. Squares are averages for the three prime WDs with their 3σ error in the mean,
while the red crosses for the cooler WD GRW+70 5824 demonstrate good agreement.
Following DeM, a smooth quadratic fit to the broadband filters (filled squares) is made,
ignoring the more deviant medium and narrow band filters (open squares). The fitted values
correct the QE as a function of wavelength, while the remaining residuals define errors in
average filter transmission. Values for F220W and F250W may need revision after on-orbit
L-flats are produced.
These residual corrections to average filter transmissions need not be the same for WFC and
HRC, because the stars fall on different positions of the physical filter; and the two flat fields
are normalized arbitrarily and not to each other. Table 8 includes the average of WFC and
HRC pivot wavelengths used for the fitting, the smooth fit values in the columns labeled "QE
Change" and residual error labeled "Filter Resid." for each filter. Figure 6 shows the final
QE, including both the long wavelength adjustments from section 5.1 and the smooth fit for
WFC. The new QE curves will be installed in the STScI Synphot software package, where
the new WFC QE curve is appropriate for the WFC operating temperature of -77C before
2006 July 4. The recently delivered WFC QE curve (Mack et al. 2007) for the post 2006 July
4 era at -80C must be modified to account for the QE changes described here for the pre 2006
July 4 epoch.
Synphot does not permit different filter curves for WFC and HRC, because the same physical
filter is used for both. Traditionally, the average of the last two columns of Table 8 has been
used to update the Synphot filter component throughput. However, the WFC/HRC
differences are often greater than the 3σ uncertainties appearing in Figures 11-12; and
therefore, the Synphot package should be modified to accept different multiplication factors
for the WFC and HRC filter transmission curves. Alternatively, existing flat field
observations with the internal lamp could be analyzed to measure the proper relative
normalization of the WFC to the HRC flat fields.
Another estimate of the typical uncertainties in these corrections can be estimated by the
average absolute difference of 0.3% between the red crosses and filled squares for the total of
19 broadband filter measurements in both cameras. Larger differences for the narrow band
filters can be caused by noise in the minimal STIS observations of the standard star spectrum
of GRW+70 5824 or by the finite resolution at H-alpha that would affect the predicted count
rate P for F658N and F660N. The models are used for the SEDs of the three prime WDs; and
the development of a high fidelity model SED for GRW+70 5824 could alleviate these
shortcomings.
23
Table 8. QE and Filter Transmission Updates
Filter
Pivot-WL
F220W
F250W
F330W
F344N
F435W
F475W
F502N
F550M
F555W
F606W
F625W
F658N
F660N
F775W
F814W
F892N
F850LP
2255
2716
3363
3500
4314
4760
5022
5580
5357
5903
6303
6520
6650
7679
8087
8914
9100
QE-Change
WFC
HRC
...
...
...
...
1.010
1.013
1.014
1.017
1.016
1.018
1.019
1.020
1.020
1.022
1.023
1.023
1.023
Filter-Resid.
WFC
HRC
1.005
1.005
1.005
1.005
1.005
1.005
1.005
1.004
1.004
1.004
1.003
1.003
1.003
1.002
1.001
1.000
0.999
...
...
...
...
1.000
0.999
1.019
0.974
0.998
1.006
0.997
1.003
1.013
0.994
0.997
1.017
1.007
0.998
1.017
0.996
0.959
0.995
1.003
0.994
1.021
0.999
1.001
1.003
0.995
1.003
0.999
0.990
0.985
1.007
6.1 After Iteration
After making the corrections for the side 2 temperatures, the EE, the changes with time, the
QE, and the filter transmissions, the average photometry for the three prime WDs is corrected
to exactly agree with the average synthetic photometry from the pure hydrogen model SEDs
for these WDs. Because the smooth QE corrections described above may have a slightly
different effect on the synthetic photometry for the cooler stars, the residuals for these five
cooler stars with respect to the WDs in Tables 4 and 6 are recomputed with the revised QE
curves. Changes are negligible with one change to the Table 4 or 6 values as big as 0.3% for
each of WFC and HRC.
7. Future Calibration Planning
If ACS and STIS are revived during SM4, then this investigation of the filter throughputs
could be continued. Mainly, more observations of the five cooler stars with both instruments
are required to refine and verify the current sparse data set and to further investigate stellar
variability. In order to verify the applicability of the ACS flat fields for red stars, a good test
for filter transmission variations with stellar spectral type would be to step VB8 around the
WFC field for the problematic filters. With hindsight and for the suggestion box of future
missions, monochromatic flats per the lab calibration campaign (Bohlin, Hartig, and Martel
24
2001) should have been obtained to assess filter uniformity near the filter cutoff wavelengths.
The few existing monochromatic flats often differ significantly in their spatial structure from
the corresponding white light flats, which suggests the presence of spatial variation in the
transmission curves for stars with extreme SEDs.
F. Boffi utilized the Synphot software to produce the current filter system throughput curves,
the pivot wavelengths, and the detector QE functions.
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