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Instrument Science Report WFC3 2014-15

Enabling Observations of Bright Stars with WFC3 IR Grisms
Susana Deustua, Ralph Bohlin and John MacKenty 30 June 2014

A

BSTRACT

The re-enabled scan mode of HST allows scan rates up to ~7.8 arcsec/sec under gyro control, providing a fast enough rate to obtain a high number of photons per pixel without saturating the detector for imaging and spectroscopy. We have demonstrated that point sources as bright as V~0 mag can be successfully observed with the WFC3/IR grisms. Program 12336 obtained scanned spectral data of bright stars with visible magnitudes between V=0 and V=7. Here we show the results from observations of Vega, the quintessential absolute flux calibrator in Astronomy. Spectra of Vega are obtained using spatial scanning in the -1st order with the G102 and G141 grisms. With this capability, the WFC3/IR channel achieves a dynamic range of over 26 magnitudes in brightness.

Introduction
This report describes the results of the first observations of Vega using the WFC3 IR grism, G141, in scanning mode. These data were obtained in the first visit of program 12336, entitled Scan Enabled Photometry. Principal motivation for this experiment is the goal of placing the WFC3 photometry directly on the Vega photometric scale. Current questions in cosmology (e.g. the nature of dark energy, the value of the Hubble constant) and astrophysics (e.g. stellar evolution, exoplanet characteristics) require improvements in the accuracy and precision of photometric measurements. For example, determining whether there is a time variable component of dark energy requires that the


uncertainty in cross-band photometry1 be < 1% for those experiments that use Supernovae. At present, the accuracy of absolute photometry is limited primarily by number of astronomical flux standards that have SI-traceable2 calibrations and wavelength/frequency ranges in which this calibration is valid. Therefore, we rely on extrapolate flux calibrations to other wavelengths

Typa Ia the small the few models to

Vega, a bright, A0V star, is the traditional absolute flux calibrator in astronomy, and, is one of only a few stars calibrated against an SI-traceable, radiometric standard. A more detailed discussion of the history of Vega's and Sirius' flux calibration can be found in Deustua, Kent and Smith (2012). The HST set of standard stars, observed by every HST instrument to determine zeropoints, derives its photometric absolute flux calibration from white dwarf (WD) models normalized to the Megessier (1995) determination of Vega's flux at 5556 å. Photometry at wavelengths in the IR, are calculated from zeropoints obtained by extrapolating the WD models to the longer wavelengths. Implicit is the assumption that because WD models are relatively simple, an excellent solution in one wavelength region is extendable to another. Thus, the zeropoints of HST's two infrared instruments, NICMOS and WFC3/IR, are calculated from the canonical HST standards: the three hot white dwarf standards GD153, GD71 and G191B2B and the solar analog star P330E ( GSC 02581-02323). In 2011, HST implemented a drift scan observing mode, thereby enabling high signal to noise photometry and spectroscopy of objects that would otherwise saturate in the traditional `stare' modeIn scan mode, HST sweeps across the field of view at a user selected rate and orientation. Motivation for making this mode available stems from exoplanet studies where the planet's star is often brighter than V=12 mag. Principal advantages of scanning are access to brighter targets and higher photometric precision as observations extend over large numbers of pixels, reducing errors arising from flatfields, cosmic rays, atypical pixels, and so forth. With spatial scanning, WFC3/IR spectroscopy of bright sources that would otherwise saturate in stare mode becomes feasible. The bright limit for stare mode grism spectroscopy is J~8.5 mag for the G102 grism and J~9 mag for the G141 grism, and, for full frame, stare mode imaging is ~13 mag. The IR subarrays have brighter limits because exposure times can be shorter than 2.93 seconds. Spectra in the -1st orders, are fainter by an average of ~ 4.68 magnitudes (a factor of 0.015) compared to +1st order. Combined with the capability of setting the scan rates as fast as 7.8 arcsec/second enables spectroscopy of very bright sources like Vega, and possibly even Sirius. At this fastest rate, the target's effective exposure time per pixel during the scan can be estimated, a priori, as: To estimate the PSF, we use the aperture correction with wavelength bin as given in Kuntschner et al (2009a) and Kuntschner et al (2009b) for the +1 orders.

1

Also referred as absolute color calibration ­ where the flux ratio in any two bandpasses is known accurately and precisely. Knowledge of the slope, but not the absolute flux, is necessary. 2 The International System of Units (SI, from the French for SystÕme International d'UnitÈs) provides ideal measurement standards for radiometry as they can be realized anywhere, at any time in the future, and are traceable to the fundamental physical units.

2


t

effective

pixel scale â PSF scan rate 0.129 "/ pix = â 2.4 7.8 "/ sec = ~ 0.04 sec

This ISR provides the first set of results for the spectral scans of Vega. Subsequent ISRs will address in further detail calibration of the sensitivity functions of the -1st orders, spatial stability, and new wavelength calibrations, with a planned summary paper that will provide the absolute photometric pedigree of WFC3 with respect to Vega.

Program 12336
This program tests the effects of different scan rates, look for variation in spatial response with position (along the columns) and brightness of measured spectrophotometry of Vega and four secondary flux standards. Target stars and their photometric characteristics are given in Table 1. Spectra of P330E, a 12th magnitude G star are taken in both +1st and -1st orders of the G141 grism to provide a cross check on the predicted sensitivities of these orders.
Table 1. Photometric Characteristics of Target Stars for Program 12336 Sp. Type Lyr (Vega) HR 7018 HD 159222 P330E* HD 93521 A0V A0V G1V G2V O9Vp B 0.03 5.690 7.21 12.972 6.79 V 0.03 5.740 6.56 12.917 7.03 R 0.1 6.1 12.564 I 0.2 5.8 12.212 J -0.18 5.725 5.342 11.769 7.499 H -0.03 5.773 5.076 11.454 7.647 K 0.13 5.753 4.998 11.379 7.696 References 2, 4 2, 5 2, 4 1 3, 5, 6

* GSC 02581-02323 is its primariy identifier in the SIMBAD database [1] Casagrande, L, Portinari, L., & Flynn, C., 2006, MNRAS, 373, 13 [2] Cutri, R.M., et al., 2MASS All Sky Catalog of Point Sources, 2003yCat.2246.0 [3] HÜg,E., et al., 2000, A&A, 355L, 27 [4] Monet, D.G., et al, 2003, AJ, 125, 984 [5] Oja, T., 1991, A&AS, 89, 415 [6] Reed, B.C., 2003, AJ, 125, 2531

Spectral scans of Vega with WFC3/IR
In this ISR we concentrate on the results from Visit 01 of Program 12336, which demonstrates the feasibility of obtaining data of bright sources. Subsequent reports will discuss

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the photometric analyses of the program. Each grism has two orbits ­ one at the fastest scan rate available, 7.8 arcseconds/second, and one at a slower rate. Two rates were chosen to test that there are no differences in photometry with scan rate. The slower scan rate was selected so that the number of electrons per pixel in the -1st order is less than ~50,000 electrons or 2/3 of the saturation level of 70,000 electrons. To measure the effect of self-persistence, and to position the star on the detector, a dark frame is taken before each pair of scans. Vega was placed at different locations on the array using POSTARGS, with the first position putting the zero order image on the edge of the detector. Subsequent POSTARGS moved Vega's zero order spectrum to the left, to mitigate against persistence and to test spatial repeatability. Table 2 lists the visit number, grism, scan rate and observing sequence. In each visit forward and reverse scans were employed. In the case of a forward scan, the POSTARG determines the starting point, whereas in a reverse scan, the POSTARG sets the end point of the scan..
Table 2. Scan Sequences for the Vega Visits Visit 1 2 3 4 Grism G141 G141 G102 G102 Scan Rate arcsec/sec 4.9 7.8 4.1 6.5 Dark Dark Dark Dark Pos. 1 scans Pos. 1 scans Pos. 1 scans Pos. 1 scans Observing Sequence Dark Dark Dark Dark Pos 2 scans Pos 2 scans Pos 2 scans Pos 2 scans Dark Dark Dark Dark Pos. 3 scans Pos. 3 scans Pos. 3 scans Pos. 3 scans

In each Vega orbit, the -1st order spectrum is placed in three separate positions on the Focal Plane Array as show in Figure 1. The first pair of observations was taken with the -1st order placed as far to the right as possible (Figure 1, leftmost image). Each successive pair of scans moves the -1st order to the left. For the first scan, the +1st order is off the detector, in the second, the zeroth order is near the middle, and in the third scan the +2nd order appears.

-2

-1

0

+1

-2

-1

0

+1

-1

0

+1

+2
nd

Figure 1. Positions on the detector of the spectral scans. Grism orders from right to left are +2 , st nd th st nd. st rd +1st, 0th, -1 and -2 . Left: 0 , -1 , -2 Slivers of the +1 on the right edge, and -3 on the left st th st nd nd st th st edge are visible. Center: +1 , 0 , -1 and -2 orders. Right: +2 , +1 , 0 and -1 orders are th visible. The 0 , +1st and +2nd orders are saturated

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Ten scans of Vega were acquired with G141 grism at a scan rate of 4.9 arcsec/second, and 12 scans at 7.2 arcsec/seconds. With the G102 filter, 7 scans were obtained at 4.1 arcsec/second, and 12 at 6.5 arcseconds. The exposure per pixel is set by the scan rate, not the sample sequence. Table 3 lists the dataset name, target, start date and time of each scan, POSTARG1, POSTARG2, and the direction of the scan. An F in this last column means a forward scan, R is a reverse scan and 0 indicates no scan. The data were reduced by the HST CALWF3 pipeline, which carries out dark subtraction, applies the linearity corrections and up-the-ramp fitting on the raw files. Each WFC3/IR raw.fits file contains N+1 extensions, where N is the number of read samples. The resulting flt.fits files contains 5 extensions, one of which is the science data. Readers are directed to the WFC3 Data Handbook for a more detailed discussion of IR data reductions. For normal spectroscopy taken in the traditional stare mode, the *flt.fits files are used for subsequent analyses, for example with the aXe spectral software package. However, when scan mode is used, the spectra is not localized to a particular part of the detector. Each read contains a portion of the scanned spectrum. Therefore the analysis for this program uses the *ima.fits files which are not flatfielded, flux calibrated, or had up-the-ramp fitting applied. However, they have been corrected for the gain and the units converted to electron per second. Specifically we use the last sample, which contains all the scanned pixels, and, from which the zeroth read (effectively the bias level) has been subtracted. The scanned spectral data were analyzed using the aXe Spectral Extraction and Visualization software for slitless spectra (Kummel et al 2009), provided as part of the STSDAS software package, in PyRAF. Because aXe was designed to be used with traditional, static spectra, we developed work arounds to the standard procedures. For example, aXe assumes that a direct image of the scene is taken at the same position just prior to the spectral image (see Kummel et al 2011) For stars brighter than ~12 mag, a direct image, regardless of filter used, will saturate a full frame exposure even for NSAMP=1. An additional complication is that the spectra are scanned, so a single direct image is of limited use. Therefore, instead of the traditional direct image, a pseudo catalog of direct image coordinates along the scan direction was generated from the measured start and end positions of the observed orders. Calculated separations of the grism orders with respect to the direct image are given by Petro (2011). The catalog contained pseudo stellar positions separated by 10 pixels in the direction of the scan (i.e. up the rows, along the columns). Once the catalog was created, the normal aXe reductions (flatfield, sky subtraction and wavelength calibration, see Table 4 for the specific reference files) were applied to the -1st order using the default aXe parameters but with an extraction width of +/- 5pixels. The resulting extracted spectra (not flux calibrated) along the scan direction were then coadded. These data were also analyzed using custom IDL procedures that instead relied on the centroid position of the 0th order in the dispersion direction for navigation to the -1st order. The same reductions were applied as in the aXe approach, and the unfluxed spectra coadded.

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Table 3. List of Observations for Visit 01, 02, 03, 04
Data Set ibtw01q5q ibtw01q6q ibtw01q7q ibtw01q8q ibtw01q9q ibtw01qcq ibtw01qdq ibtw01qeq ibtw01qfq ibtw01qgq ibtw01qiq ibtw01qjq ibtw01qkq ibtw02ghq ibtw02giq ibtw02gjq ibtw02gkq ibtw02glq ibtw02gmq ibtw02goq ibtw02gpq ibtw02gqq ibtw02grq ibtw02gsq ibtw02gtq ibtw02gvq ibtw02gwq ibtw02gxq ibtw03jkq ibtw03jlq ibtw03jmq ibtw03joq ibtw03jpq ibtw03jqq ibtw03jsq ibtw03jtq ibtw03juq ibtw04a1q ibtw04a6q ibtw04a7q ibtw04a8q ibtw04a9q ibtw04aaq ibtw04acq ibtw04adq ibtw04aeq ibtw04afq ibtw04agq ibtw04aiq ibtw04ajq ibtw04akq ibtw04alq V V V V Target DARK ALF-LYR ALF-LYR ALF-LYR ALF-LYR DARK ALF-LYR ALF-LYR ALF-LYR DARK ALF-LYR ALF-LYR ALF-LYR DARK ALF-LYR ALF-LYR ALF-LYR ALF-LYR DARK ALF-LYR ALF-LYR ALF-LYR ALF-LYR DARK ALF-LYR ALF-LYR ALF-LYR ALF-LYR ALF-LYR ALF-LYR DARK ALF-LYR ALF-LYR ALF-LYR DARK ALF-LYR ALF-LYR DARK ALF-LYR ALF-LYR ALF-LYR ALF-LYR DARK ALF-LYR ALF-LYR ALF-LYR ALF-LYR DARK ALF-LYR ALF-LYR ALF-LYR ALF-LYR Date 2011 2011 2011 2011 2011 2011 2011 2011 2011 2011 2011 2011 2011 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 2012 12 12 12 12 12 12 12 12 12 12 12 12 12 08 08 08 08 08 08 08 08 08 08 08 08 08 08 08 11 11 11 11 11 11 11 11 11 11 11 11 11 11 11 11 11 11 11 11 11 11 11 11 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 14 20 20 20 20 20 20 20 20 20 26 26 26 26 26 26 26 26 26 26 26 26 26 26 26 Start Time 05:30:01 05:31:53 05:34:26 05:35:46 05:37:06 05:50:27 05:51:58 05:53:18 05:54:38 05:57:04 06:07:56 06:09:16 06:10:36 02:57:04 02:59:25 03:00:45 03:02:05 03:03:25 03:05:36 03:17:02 03:18:22 03:19:42 03:21:02 03:23:13 03:24:32 04:24:50 04:26:10 04:27:30 12:39:45 12:42:06 12:44:25 12:56:01 12:58:33 13:01:05 13:13:31 13:15:00 13:16:24 00:35:27 00:59:49 01:01:06 01:02:23 01:03:40 01:05:52 01:18:30 01:19:47 01:21:04 01:22:21 01:24:33 01:35:36 01:36:53 01:38:10 01:39:27 POSTARG1 arcsec 62.0 66.0 62.0 62.0 62.0 42.5 42.0 42.0 42.0 15.0 15.0 15.0 15.0 0.0 66.0 66.0 66.0 66.0 0.0 42.0 42.0 42.0 42.0 0.0 15.0 15.0 15.0 15.0 66.0 66.0 0.0 46.0 46.0 46.0 0.0 26.0 26.0 0.0 66.0 66.0 66.0 66.0 0.0 46.0 46.0 46.0 46.0 0.0 26.0 26.0 26.0 26.0 POSTARG2 arcsec -62.0 149.919 149.919 -117.233 149.919 -63.0 149.919 -117.233 149.919 -63.0 149.919 -117.233 149.919 0.0 209.170 -159.598 209.170 -159.598 0.0 209.170 -159.598 209.170 -159.598 0.0 212.170 -156.598 212.170 -156.598 152.619 -86.509 0.0 141.854 141.854 141.854 0.0 141.854 -86.509 0.0 190.284 -127.859 190.284 -127.859 0.0 190.284 -127.859 190.284 -127.859 0.0 190.284 -127.859 190.284 -127.859 Scan Direction 0 F F R F 0 F R F 0 F R F 0 F R F R 0 F R F R F R F R F R F F F 0 F R 0 F R F R 0 F R F R 0 F R F R

VVVVVVV V V V V V V V V V V V V V -

VVVVVV V V V V V V V V V V V -

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Table 4. GRISM Reference Files configuration file: sensitivity -1st order: flat-field cube: master sky image: WFC3.IR.G141.V2.5.conf WFC3.IR.G141.m1st.sens.2.5.fits WFC3.IR.G141.flat.2.fits WFC3.IR.G141.sky.V1.0.fits WFC3.IR.G102.V2.0.conf WFC3.IR.G102.m1st.sens.2.fits WFC3.IR.G102.flat.2.fits WFC3.IR.G102.sky.V1.0.fits

Results
The raw counts per pixel at the slow scan rates are measured to be ~50000 electron/pixel at the slow scan rates for both the G102 and G141 grisms, and between 30000 ­ 35000 electrons/pixel at the faster scan rates Saturation in the WFC3/IR detector is set at ~70000 electrons, where the . Signal to noise per pixel averages SN=200, and exceeds SN=3000 for each coadded spectrum per scan. In table 5, we provide the measured count rate in the ima file and the appropiate count rate based on the `onsource' time of each pixel. We also provide, for reference, the equivalent quantities for the brightest pixel. Because the HST/WFC3 pipeline is designed for `stare' mode observation, the count provided in the *ima.fits are incorrect. CALWF3 divides each pixel by the exposure contained in EXPTIME keyword in the image header. During a scan, the exposure time by the scan rate and not the sample sequence, therefore the correct count rate per pixel has recalculated: observed count rate = calwf3 count rate * exposure time * scan rate / pixel scale Thus, the count rate in the *ima.fits files underestimates the true number of electrons per second per pixel. From the values listed in Table 5 we calculate that the count rates differ by 1.6% between the slow and fast G141 scans, and by 3.4% for the G102 scans.
Table 5: Measured Counts from the Vega Spectral Scans G141 Commanded scan rate: (arcsec per sec) Total exposure time in the scan (seconds) calwf3 count rate in ima.fits file (electrons per second) Average on source exposure time per pixel (seconds) Average number of electrons per pixel of the scanned spectrum in ima.fits files "true" count rate per pixel of the scanned spectrum in ima.fits file (electrons per second) Number of electrons in the brighest pixel in the scanned spectrum "true" count rate for the brightest pixel in the scanned spectrum (electrons per second) 4.9 20.526 1 564 0.027 32 103 1 210 026 47 200 1 779 450 7.2 11.729 1 832 0.018 21 488 1 190 095 32 800 1 818 923 G102 4.1 26.391 1 167 0.032 30 789 971 039 63 700 2 008 942 6.5 14.661 1 371 0.020 20 097 1 004 860 40 100 2 003 763

rates time is set to be

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Flux (electrons/sec)

6·10 4·10 2·10

8

8

F8gure 2 shows the coadded spectra for the two grisms. The saddle i response of that grism's -1st order; the peak G102 response peaks at wavelength values indicate a minus order. Comparing the slow to the 0 G102 shows little difference between the two. However, the same is 1.0 1.2 1.4 1.6 1.8 cause is under investigation, but obvious suspects are persistence or flat
Wavelength (microns)

shape in G141 is the real ~0.8 micros. Negative fast scanned spectra for not true for G141. The field effects.

6·10
Flux (electrons/sec)

8

Vega: coadded spectra G102
8·10
Flux (electrons/sec)

8

Vega: coadded spectra G141

4·10

8

6·10 4·10 2·10

8

8

2·10

8

8

0 0.5

0
0.6 0.7 0.8 0.9 1.0 1.1 Wavelength (microns) 1.2

1.0

1.2 1.4 1.6 Wavelength (microns)
st

1.8

Flux (electrons/sec)

Figure 3 Coadded unfluxed spectra of Vega at each scan ratVega: coadded spectra G102 d e for the -1 orders of the G141 (top) an G102 (bottom) WFC3 IR grisms. Wavelength scale i6·10microns, negative numbers indicate negative s in 8 order. Solid lines are the slow scan rates; dashed lines are the fast scans. Top: Paschen is the dip at 1.28 microns, and the series of features between 1.6 and 1.7 microns are Br 13,12, 11. Bottom: Pa is just visible at 1.09 microns, and Pa at 1.05 microns.

4·10

8

G1412·108 ibtw01

6·10

8

DN/sec for V-ALF-LYR

0 0.5
8

0.6

0.7 0.8 0.9 1.0 1.1 Wavelength (microns)

1.2

4·10

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0 -1.8
wfc_coadd 2-Jan-2013 23:14:48.00

-1.6

-1.4 -1.2 -1.0 Wavelength (microns)

-0.8

-0.6

Figure 3. Visit 01; slow scan, -1 order spectra in G141 coadded at each scan position, after dark subtraction, flatfielding, dispersion correction and rectification. Note the 10% variation in the counts rates at wavelengths longer than 1.2 microns (left side of the plot). This is likely due to flat field variations. 8

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Figure 3 is a plot of the scanned spectrum (coadded along the scan direction) for each position on the detector. Plotted is the count rate in DN/sec versus wavelength. At wavelengths longer than 1.1 microns, the spectra at different detector positions differ from each other by up to 10%. However, we find that the reproducibility of these spectra is better without applying the flatfield. This implies that the flat fields are a source of significant error, and may not be flat across the detector. As a check on wavelength calibration uncertainites, we compared a G141 -1st order extracted spectrum of Vega to a 1st order, G141, spectrum of the white dwarf GD71. Plotted in Figure 4 are the +1st order spectrum of GD 71 (dotted curve) and a -1st order spectrum of Vega (solid curve). Vega's spectrum is not corrected for sensitivity response. A vertical, dotted line marks the Pa line at 1.282 µm. The fit is quite good but perhaps with a small shift redward of the -1st order relative to the +1st order. Re-determining the wavelength solutions for the grisms for the minus orders is now under way. From these observations we calculate the brightest limiting magnitude to obtain ~35900 source electrons per pixel at several scan rates for both the +1st and the -1st grism orders, after subtracting the dark frame and the zeroth read. The equivalent pixel will have ~50 000 e- in the raw file. In Table 6 we list the limiting magnitude (in electrons/second) of a point source for each grism that meets this condition for various scan rates. For reference, the ETC uses a count rate of 1.57 x 1010 electrons/sec for Vega, from the model fit to the STIS spectrum obtained by Bohlin & Gilliland (2004). Sirius, the brightest star in the sky and also an SI-traceable flux standard, has J=-1.391 mag, about 40% brighter than the limit of J=-1.03 mag listed for G102 at 7.8 arcsec/sec. At the fastest scan rate, the expected counts corresponds to 62000 electrons, just below the nominal saturation limit of 70000 electrons. In principle then, Sirius is observable.

Figure 4. Comparison of the -1 order spectrum in G141 of Vega (solid line) to the 1 order spectrum of the white dw arf GD71 (dotted line). The vertical line marks the position of the Pa line at 1.282 microns. st st Note the redward shift of the -1 order relative to the +1 order.

st

st

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Table 6. Brightest limiting magnitudes for each grism order as a function of scan rate G141 Grism Scan Rate arcsec/sec 0 0.1 1 3 5 7.2 Exposure time sec/pixel 2.93 1.30 0.13 0.043 0.026 0.018 Average Count Rate (per pixel) e-/sec 1.18 x 10 2.66 x 10 2.66 x 10 7.98 x 10 1.33 x 10 1.91 x 10
1

G102 Grism Average Count Rate (per pixel) e-/sec 1.23 x 10 2.76 x 10 2.76 x 10 8.29 x 10 1.38 x 10 1.99 x 10
1

Limiting Brightness J m ag2 -1st order +1st order 9.52 8.64 6.14 4.95 4.39 4.00

Limiting Brightness J m ag2 -1st order +1st order 9.70 8.81 6.31 5.12 4.57 4.17

4 4 5 5 6 6

4.84 3.96 1.46 0.27 -0.29 -0.68

4 4 5 5 6 6

4.59 3.70 1.20 0.01 -0.54 -0.94

7.8 0.017 2.07 x 10 6 -0.77 3.91 2.16 x 10 6 -1.03 4.08 1 The count rate needed to obtain 35900 electrons per pixel at given scan rate (source only, not including dark and zeroth read subtraction which add ~11000 electrons), or ~50000 electron per pixel in the raw frame. 2 Average magnitude difference between the +1st and -1st orders is ~4.68 in G141, and ~5.11 in G102

Future Work Work under way is to to improve the wavelength calibration of the -1st orders across the detector, check and possibly improve the flat field calibration, and investigate the difference in photometry between slow and fast scans.

Conclusions
Based on this analysis we conclude that a) Bright point sources, as bright as J=-1 mag can be successfully observed with the HST WFC3/IR channel using spatial scanning and exploiting the lower sensitivity of the -1st grism orders. This gives an almost 26 magnitude dynamic range for the WFC3/IR channel. b) Additional effort is being placed on calibrating the -1st orders to the same level of confidence as for the +1st and +2nd orders We note that at the time WFC3 was installed in the Hubble Space Telescope in 2009, there was no drift scan mode available, and no expectation that IR grism orders other than the +1st and +2nd would be used.

Acknowledgements
We thank Merle Reinhardt for his help in putting together the observing program, Peter McCullough for generously letting us piggy-back off his exoplanet IDL code and for many useful discussions on understanding scanned spectra, Nor Pirzkal for pointers on the subtleties of grism spectroscropy, and, Adam Riess for a careful reading of the manuscript.

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References
R. Bohlin & R. Gilliland, Hubble Space Telescope Absolute Spectrophotometry of Vega from the Far-Ultraviolet to the Infrared, 2004, AJ, 127, 3508 Deustua, S., Kent, S., and Smith, J.A. (2012). Absolute Calibration of Astronomical Flux Standards, Planets, Stars and Stellar Systems, Eds. Oswalt, Terry D.; Bond, Howard E., ISBN 978-94-007-5617-5. Springer Science+Business Media Dordrecht, 2013, p. 375 Hayes, D. S.; Latham, D. W., A rediscussion of the atmospheric extinction and the absolute spectral-energy distribution of VEGA, 1975, ApJ, 197.593H KÝmmel, M.; Walsh, J. R.; Pirzkal, N.; Kuntschner, H.; Pasquali, A. The Slitless Spectroscopy Data Extraction Software, 2009, PASP,121, 59 M. KÝmmel, J. Walsh, H. Kuntschner and H. Bushouse, aXe User Manual version 2.3, 2011, http://axe.stsci.edu/axe/manual/html/index.html L. Petro, , IR Grism Subarray Design and Use , WFC3 Technical Instrument Report 2010-03

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