Appendix H. List of Figures
Figure 2.1: IRAC Cryogenic Assembly model, with the top cover removed to show the inner components.
Figure 2.2: IRAC optical layout, top view. The layout was similar for both pairs of channels; the light entered the doublet and the long wavelength passes through the beamsplitter to the Si:As detector (channels 3 and 4), and the short wavelength light was reflected to the InSb detector (channels 1 and 2).
Figure 2.3: IRAC optics, side view. The Si:As detectors are shown at the far right of the figure. The InSb arrays are behind the beamsplitters.
Figure 2.4: Spectral response curves for all four IRAC channels. The full array average curve is displayed in black. The subarray average curve is in green. The extrema of the full array per-pixel transmission curves are also shown for reference. The red curves are for the pixel in the array with the highest nominal wavelength bandpass and the blue curves for the pixel in the array with the lowest nominal wavelength bandpass. For definition of the “nominal wavelength” see Hora et al. (2008). The spectral response curve data are available at https://irsa.ipac.caltech.edu/data/SPITZER/docs/irac/calibrationfiles/spectralresponse/ .
Figure 2.5:Optical image distortion in IRAC channels in the cryogenic mission (above) and warm mission (below). The panels show the image distortions as calculated from a quadratic polynomial model (above) and a fifth-order polynomial (below) that was fit to in-flight data. The magnitude of the distortion and the direction to which objects have moved from their ideal tangential plane projected positions is shown with arrows. The length of the arrows has been increased by a factor of ten for clarity. The maximum positional deviations across the arrays for this quadratic distortion model are less than 1.3, 1.6, 1.4, and 2.2 pixels for channels 1 - 4, respectively. The derivation of the pixel scales that are listed in Table 2.1 fully accounted for the quadratic (cryogenic data) and fifth order (warm data) distortion effects shown here.
Figure 2.6: Cryogenic linearity curves for the IRAC detectors. As noted in text, the warm detectors were roughly linear to 30000 DN.
Figure 2.7: Fowler sampling times for one pixel (Fowler N=4). The Pn (n=1,2,3,4) show the “Pedestal” readouts, and the Sn show the “Signal” readouts. Tex is the effective exposure time, and Tf – Ti is the “frame time,” or total time to obtain one IRAC image. The reset part of the sketch is not at the same time and voltage scale as the rest of the figure.
Figure 2.8: IRAC point source sensitivity as a function of frame time, for low background. The cryogenic values are on the left, and the warm mission values are on the right. To convert to MJy/sr, see equation (2.13). Subarrays are plotted in blue.
Figure 2.9: Similar to Figure 2.8, but for medium background.
Figure 2.10: Similar to Figure 2.8, but for high background.
Figure 3.1: IRAC AOT template. The options for selecting observing mode, fields of view, frame repeats, dithering and mapping are shown.
Figure 3.2: IRAC dither patterns for the “large” scale factor. The x- and y-axes are in pixels.
Figure 3.3: Characteristics of the cycling dither pattern. The x- and y-axes are in pixels in the left column plot. The histogram x-axis (separation) is in units of pixels (middle and right plots).
Figure 4.1: IRAC instrument skydark images. These images were taken during a normal campaign, and using 100 second frame time skydarks.
Figure 4.2: Row pixel values (in DN) of a shutter dark and a regular skydark. This plot shows the channel 2 skydark (red) pixel values along a row compared to the pixel values in a channel 2 shutter closed dark (black) in the same row. These observations used the 100 second frame times. The difference between the measurements is consistent with the astronomical background at the skydark position (1 DN = 3.7 electrons).
Figure 4.3: Nominal 30 second frame time skydarks (first row) and 30 second frame time shutter darks (second row) for channels 1 and 2 with relative scaling. The difference between the skydark and the shutter dark (last row) is also displayed.
Figure 4.4: IRAC superskyflats. Shown are the cryogenic mission (top two rows, including channels 1 – 4) and the warm mission (bottom, including channels 1 and 2) superskyflats.
Figure 4.5: Array location-dependent photometric correction images. The cryogenic mission corrections are at the top, and the warm mission corrections at the bottom. For the cryogenic mission correction images, channel 1 is in the upper left, channel 2 in the upper right, channel 3 in the lower left, and channel 4 in the lower right. White is the largest value (about 1.046) and black is the smallest value (about 0.915). For the warm mission correction images, channel 1 is on the left, and channel 2 on the right. Again, white is the largest value (about 1.01757, 1.03167 in channels 1 and channel 2), and black is the smallest value (0.9476, 0.9223).
Figure 4.6: IRAC cryogenic pixel response model, showing intrapixel sensitivity variations as a function of pixel phase. Only channel 1 (3.6 μm) and channel 2 (4.5 μm) have significant variations.
Figure 4.7: Distribution of normalized corrected calibration star flux densities. There is a difference in the absolute photometry of stars taken in different observing modes (shown here in the different colored distributions). The legend lists the different observing modes (full array vs. subarray, and long exposure times vs. short exposure times) as well as the number of photometry points per distribution. Channel 1 is on the left, channel 2 is on the right.
Figure 4.8: Cleveland dot plot showing medians and one sigma uncertainties of the distributions of photometry from four different modes shown in Figure 4.7 (subarray vs. full array and staring vs. dithering). The names of the three calibration stars considered here are shown on the y-axis, and the data points for different modes of photometry for each star are offset slightly in the y-direction for ease of viewing. Staring mode observations in the subarray and full array modes were taken at the same locations on the array (on the sweet spot pixel).
Figure 4.9: Images of the IRAC point response function (PRF) tables. These are shown at 3.6 (upper left), 4.5 (middle, top row), 5.8 (left, bottom row), and 8.0 µm (middle, bottom row) from the cryogenic mission and from the warm mission at 3.6 (right, upper row) and 4.5 µm (right, bottom row). The PRFs were generated from models refined with in-flight calibration test data involving a bright calibration star observed at several epochs. The PRFs are shown as they appear with 1/5 the native IRAC pixel sampling of 1.2 arcseconds to highlight the core structure. The PRFs were calculated using the Simfit routine in Hoffmann et al. (2004). The images are 128 × 128 1/5 native IRAC pixels in size (about 31 arcseconds × 31 arcseconds).
Figure 4.10: “Best” (left) and “worst” (right) channel 1 PRF images from the set of PRFs available for the cryogenic mission. A logarithmic stretch was used to highlight the low-level emission. The core of the “worst” image is a larger than the “best” image, and it has a ≈ 60% higher noise pixel value.
Figure 4.11: Distribution of noise pixel values over the channel 1 field of view in the cryogenic mission. The image matches the BCD orientation, with pixel (1,1) in the lower left corner. The mean point source noise pixel values (averaging over many different pixel phases) were measured at the 25 positions indicated by the small squares on the image, and a quadratic surface was fit to the data, plotted here as grayscale and contours. The lowest noise pixel value is 6.79 at (x,y) position (163,128). The highest noise pixel location, in the upper left corner, is 1.6 times this, requiring 1.6 times the integration time to achieve the same signal-to-noise as for the lowest noise pixel position.
Figure 4.12: The IRAC extended PSFs. The cryogenic PSFs are shown on the left, with channel 1 PSF in upper left, channel 2 PSF in the middle (upper row), channel 3 in lower left, and channel 4 in the middle (lower row). The warm extended PSFs are shown on the right, with channel 1 on top and channel 2 at the bottom. See the text for how these were constructed.
Figure 4.13: Reaction wheel speed and centroid position as a function of time. Wheel speed changes are a few seconds in duration and this entire AOR is 11.5 hours in duration. This particular AOR shows a reaction wheel speed shift with a return to the original position. Grey lines correspond to the wheel speeds in RPM (left y-axis). Spitzer had four reaction wheels, and they all changed speed simultaneously about 75% of the way through this observation. Blue and green data correspond to the X and Y centroids of the target in the observation (right y-axis).
Figure 4.14: Two 0.02 second frame time samples taken during a long staring observation. The total monitoring period was 42,007.38 seconds or 11.67 hours. The left panel shows the data at the very start of the stare after the initial slew to the object, and the right panel shows the jitter at the very end of the monitoring period.
Figure 4.15: Spitzer pointing during the first 100 seconds after a PCS offset to the object position. The plot shows the pointing during the first 12 subarray data sets using 0.1 second frame time (6.4 seconds per subarray set). The first set, plotted in darker green, shows up on the right side of the plot, offset from the other data sets, and showing back-and-forth oscillations of about 0.05 pixels peak-to-peak. The following data sets show a much-reduced motion in successive frames, with a random walk around the reference position.
Figure 4.16: Two examples of the drift during 100-second periods of a long staring observation. The pointing seems to walk randomly over a region of ≈ 0.05 pixels in extent, sometimes dwelling at a position for some length of time, and then moving on a path around the pointing center.
Figure 4.17: Distribution of long-term drift in both the X and Y directions as a function of time on the left and as a function of nearest in time downlink pitch angle on the right. These plots include all staring mode observations taken during the warm mission. Both plots are violin plots. Each violin-shaped object in the figure shows the median value and range for all the observations in that particular year (or at that particular pitch to earth point). Dashed lines show the median and inner quartiles of the distributions. The median values are printed at their locations. They can be thought of as smoothed histograms. Color variations between the violins have no meaning other than different years (pitches) have different colors.
Figure 4.18: Short-term drift. Top: duration, length, and slope of the short-term drift as a function of the pitch angle of the observation itself. Bottom: same short-term drift parameters as a function of the change in the pitch angle from the previous to the current observation. Observations are color coded by exposure time, where shorter exposures are red and longer exposures are blue. The scale and rate of the short-term drift correlate with the pitch angle and the duration of several previous observations.
Figure 4.19: IRAC jitter observed during a long staring observation of the Galactic Center on 2013 December 10 (PID=10060). The X and Y offsets relative to the reference position are shown for both the X and Y directions on the IRAC array. Both directions show the oscillation (approximately 0.05 pixels) caused by the spacecraft heater, with about a 40-minute period. There is also an approximately linear drift in both dimensions over the length of the observation, larger in Y than in X.
Figure 4.20: Amplitude and strength of the battery heater oscillation over the 2010 – 2017 period of the warm mission. Both plots are violin plots. Each violin shaped object in the figure shows the median value and range for all the observations in that particular year. Dashed lines show the median and inner quartiles of the distributions. Color variations between violins have no meaning other than different years have different colors.
Figure 4.21: IRAC pointing variations observed during a long staring observation on 2016 July 25 (PID=12034). This observation was similar to that in Figure 4.19, but here we do not see the spacecraft heater oscillation.
Figure 4.22: The 3.6 µm (left) and 4.5 µm (right) calibration star normalized flux densities as a function of time from 2013 to 2020. All flux densities are measured by aperture photometry and include 21 primary and secondary calibrators denoted with different colors. The line (and gap) in December 2015 denotes an anomaly in which IRAC was off and no observations were taken. There is no difference seen before or after this anomaly. See the text for the meaning of the various colors, and more information about the plots.
Figure 4.23: Median value of the 12 second skydarks plotted over time since the beginning of the warm mission. The vertical dotted lines denote two anomalies in which the array was reset. The median value was observed to settle to a new nominal level after the anomalies. The red lines give the shape of the predicted seasonal zodiacal variation normalized to the median value of the skydarks in January 2012, to help with the search for any trend changes in skydark values. See the text for more information.
Figure 4.24: Average number of pixels per second affected by cosmic rays in the 100 second skydarks since the beginning of the warm mission. The spike in 2012 is due to a solar flare that occurred during the time the skydark calibration data were being taken.
Figure 4.25: Number of pixels determined to be hot, noisy, or dead during each observing campaign in the 100 second skydarks since the beginning of the mission. Hot and noisy pixels changed over time as some pixels recovered and others became hot or noisy. Dead pixels only increased by 10 - 20 pixels over the entire mission.
Figure 5.1: Data flow for processing a raw IRAC science DCE into a BCD.
Figure 5.2: Effect of insbposdom . insbposdom worked only on the two InSb arrays (channels 1 & 2) and reversed the sense of intensities. Left is before insbposdom ran and right is after.
Figure 5.3: Diagram of the wrapping of negative values due to truncation of the sign bit.
Figure 5.4: Application of iracwrapcorr to channel 1 data. The many apparently hot pixels were actually wrapped negative values, which were detected based on their data values vastly exceeding the physical saturation value for the detectors, and corrected by subtracting the appropriate value. Real hot pixels did not exceed the physical saturation value, and hence were not changed.
Figure 5.5: Illustration of bit truncation. Bit truncation was used by IRAC for ground transmission, necessitating iracnorm. The internally stored 24-bit word was truncated to 16 bits, with a sliding window set by the barrel shift value. The figure illustrates the ABARREL=4 case.
Figure 5.6: Correction of cable-induced bandwidth error by iracebwc. The illustrated data show a cosmic ray hit.
Figure 5.7: First-frame effect: dark counts as a function of interval between frames. This figure is for a 30 second exposure frame.
Figure 5.8: Before (left) and after (right) correction of pseudo-muxbleed for channel 1. Shown is a bright source within a calibration AOR and a background of sources under the muxbleed limit.
Figure 5.9: Transposition of an IRAC channel 1 dark image by the imfliprot module.
Figure 5.10: Subtraction of the median background from the readout channel images. The image on the left shows the original image divided into each of its four readout channels. The image in the middle is the median of the four readout channels. This makes the muxstripe much more apparent in the fourth readout channel image (on the right). The location of the muxstripe is shown with black arrows. 111
Figure 5.11: Profiles showing the column median versus row values for identifying muxstripe. The x-axis is in pixel units. The muxstripe is now identifiable between row pixels 125 and 200 (significantly lower values than the median background).
Figure 6.1: The orientation of the (C)BCDs on the sky.
Figure 7.1: Full array CBCD image showing rows of missing data (dark band across the image). This observation is in channel 2, 12 second HDR frame, PID=80254, AORKEY 44050944.
Figure 7.2: Erroneously inserted saturation correction stars in a CBCD image (left), compared to the original BCD image (right). These data are from PID=46, AORKEY 12484352. The arrows point to the sources erroneously inserted by the pipeline. There are column pull-down and banding correction artifacts in the CBCD image as well, inserted by overactive pipeline corrections.
Figure 7.3: Images showing the muxbleed effect (the horizontal line on both sides of a bright stellar image). The pixels on the left side of the bright source are pixels on rows following the row in which the bright source was located (and have wrapped around in the readout order of the array). The vertical (white) lines are due to the so-called “column pull-down” effect. These are 12-second BCD frames in IRAC channel 1, taken from IRAC PID=618, AORKEY 6880000. Pinstriping is also seen in these images.
Figure 7.4: Demonstration of the S18 pipeline muxbleed removal. The image on the left is before and the one on the right is after the correction. These are First Look Survey channel 1 data, taken from AORKEY 4958976. Note that the brightest star in the upper-left corner is heavily saturated and the current muxbleed scheme can correct muxbleed from a saturated source also.
Figure 7.5: A typical bandwidth effect trail in channel 4, in a 30 second frame. These data were taken from PID=1154, AORKEY 13078016.
Figure 7.6: The bandwidth effect when a bright object is in the last four columns. IRC+10216, strongly saturated, is just off the right side of the channel 3 array. Even the filter ghost is saturated. The bandwidth effect appears on the left side of the array. These data were taken from PID = 124, AORKEY 5033216.
Figure 7.7: IRAC channel 1 (left) and channel 2 (right) observations of a crowded field with column pull-down apparent from the brightest sources. Note that the brighter sources affect a larger number of columns. These data were taken from PID=613, AORKEY 6801408.
Figure 7.8: Channels 1 and 2 (top) and 3 and 4 (bottom) showing inter-channel crosstalk (dark spots near the center of the lower panels).
Figure 7.9: Median of channel 1 images from a calibration observation performed after observing Polaris. The five bright spots are persistent images from staring at the star while observing, while the set of criss-crossing lines were generated by slews between the various pointings. These observations were taken from PID=19, AORKEY 3835904.
Figure 7.10: Residual image brightness decay as a function of time interval since exposure to a first magnitude source at 3.6 μm. The residual is compared to three times the noise in the sky background as measured in an equivalent aperture. The fitted exponential decay function is plotted as the dot-dashed line. These curves have been smoothed to mitigate flux jumps due to sources at the position of the original source in subsequent images.
Figure 7.11: Residual image examples: a. channel 1 positive residual images near the center of the array, PID=90175, AORKEY 47943424; b. same as previous but showing the median image of the observation that has also positive slew residuals; c. channel 1 positive slew residuals, PID=80096, AORKEY 45585920; d. channel 2 negative residual images, PID=90109, AORKEY 47828736.
Figure 7.12: More residual image examples; a. channel 1, PID=90124, AORKEY 48337408 (the red arrow is pointing to an image residual); b. channel 2, negative slew residuals and column pull-down residuals PID=70044, AORKEY 40840192; c. channel 1 bright slew residual, PID=61009, AORKEY 35354880; d. channel 1 bright residual image from Reuleaux pattern dithering in center, PID=80015, AORKEY 42191104.
Figure 7.13: An image of the M51 system, showing an overlay of the IRAC fields of view, with the scattered light origin zones for channels 1 and 2 overlaid.
Figure 7.14: Channel 1 image showing scattered light on both sides of a bright star. The scattered light patches are marked with white “S” letters. The images were taken from PID=30 data.
Figure 7.15: Similar to Figure 7.13 but for channel 2.
Figure 7.16: Similar to Figure 7.13 but for channel 3.
Figure 7.17: Similar to Figure 7.13 but for channel 4. The scattered light patches are pointed to by black arrows.
Figure 7.18: Typical image sections showing the banding effect. These are channel 3 (left) and channel 4 (right) images of the same object (S140), adopted from a report by R. Gutermuth. These data were taken from PID=1046, AORKEY 6624768.
Figure 7.19: Filter and beamsplitter ghosts in IRAC. Top: Channel 1 and 2 mosaics of saturated images of Fomalhaut near the center of each array from PID=90, AORKEY 4875776. The panels are 180 pixels × 140 pixels of 0.6 arcseconds; the greyscale is logarithmic. Bottom: channel 3 and 4 single, highly saturated images of Sirius near the center of each array from PID=1156, AORKEY 16412416. The panels are 90 raw pixels × 70 raw pixels; the greyscale stretch is square-root. The channel 3 and 4 filter ghosts are nearly saturated and the bandwidth effect is saturated. Even the optical banding (light diffracted and internally reflected within the arrays) in the column containing the star center suffers the bandwidth effect, which appears as bright vertical spikes 4 and 8 pixels to the right of the center of the star.
Figure 7.20: Pupil ghost in channel 2 from V416 Lac.
Figure 7.21: Part of the channel 1 mosaic (from observations in PID=181; AORKEYs 5838336, 5838592, 5839872, and 5840128) of the SWIRE field near Mira showing the 24 arcminute radius ring of stray light from the telescope. The dimensions of the image are about 0.7 degrees (x-axis) by 1.1 degrees (y-axis).
Figure 7.22: Channel 2 images from the SWIRE map showing stray light splotches from Mira, which was about 30 arcminutes away. Successive pairs of images were slightly dithered. The last pair is about 5 arcminutes from the first pair, but has a similar splotch. Note the absence of any stray light in the second image, though it was centered only a few pixels away from the first image. The images are from PID=181, AORKEY 5838336; EXPID 187 - 192, 199, and 200.
Figure 7.23: The central 128 pixels × 128 pixels of IRAC 12-second images taken on January 20, 2005 during a major solar proton event. Channels 1 and 2 are top left and top right; channels 3 and 4 are bottom left and bottom right. Except for the bright star in channels 1 and 3, almost every other source in these images is a cosmic ray. These data are from observations in PID=3126.
Figure 8.1: Correcting photometric measurements of calibration star NPM1+67.0536 using the pixel phase and array location-dependent corrections. Black curves show the cumulative distribution of aperture photometry measurements throughout the cryogenic mission, for which a value is within a given per cent of the mean. Blue curves show the distribution after correction using an earlier (now obsolete) set of location-dependent functions. Red curves are the result of the improved (now final) correction. A gray horizontal line at 68% (1σ for a Gaussian distribution) is labeled with the corresponding spread in values for each of the three distributions.
Figure 8.2: Extended source flux aperture correction factors; solid lines represent exponential function fits to the data. Also indicated are correction factors derived from zodiacal light tests, and Galactic HII region tests (e.g., GLIMPSE vs. MSX).
Figure 8.3: Extended source flux correction factors for galaxies (solid lines) versus the PSF aperture correction factors (dotted lines). The main difference between the two is the truly diffuse scattering internal to the array.
Figure 8.4: Noise versus binning length in channel 1. To make this plot, the surface brightness was measured in nine regions across an object-masked mosaic. These regions are not near bright galaxies, stars, or diffuse plumes. The noise is defined as the standard deviation of those nine regions. The box size is incrementally increased until the box length is many hundreds of pixels. For reference the solid line shows the expected linear relation.
Figure 8.5: Similar to Figure 8.4 but for channel 2.
Figure 8.6: Noise as a function of exposure time (number of frames) in channel 1. The results from the warm mission data are shown with crosses and the expected behavior with the solid line. The results from the cryogenic mission are shown with open squares and the expected behavior with the dashed line.
Figure 8.7: Similar to Figure 8.6 but for channel 2.
Figure 8.8: Noise in electrons as a function of binning scale showing the relative strength of instrument systematics. Noise contributions from the source and background are shown as diagonal lines. Known instrumental systematics are shown as horizontal lines. For reference, also shown are the strengths of a hot Jupiter (1%) transit, 0.1% eclipse, and a 70 ppm super-Earth transit.
Figure 8.9: Channel 2 photometric gain versus diagonal sub-pixel position, based on a crude Gaussian model of the pixel. The vertical ranges show the amount of gain variation over about 0.2 pixels at the peak of gain (green) and near the corner of the pixel (red).
Figure 8.10: The map of intrapixel photometric gain, or “pmap” for channel 2 post-cryogenic data. This is the pmap built from data collected as of 2014/09/04. Note that gain is not symmetric around the center of the sweet spot and that the peak is not at the nominal center of the pixel.
Figure 8.11: Distribution of average noise pixels per AOR vs. exposure time. All ≈ 1300 staring mode AORs from the warm mission through June 2017 were used. Each of the gray shapes represents the distribution for the corresponding exposure time. Red lines show the median value per exposure time and are marked with that median noise pixel value. The total number of AORs per exposure time bin is printed at the top of the plot. We use 3σ clipping to remove outliers, such as those that fell far away from the sweet spot.
Figure 8.12: x, y, noise pixels, and flux values as a function of time over five simulated AORs. Note the noise pixel spikes at 16 and 30 hours which are seen in the noise pixel and flux plots, but not in the x and y position plots. This simulated dataset is of WASP 52.
Figure 8.13: Signal-to-noise ratio versus aperture size for the WASP-14b exoplanet, normalized to an aperture size of 2.5 pixels in radius.
Figure 8.14: Residual nonlinearity in warm mission channel 2 data. These plots show linearity-corrected aperture flux as a function of well depth (fluence) for a set of targets, measured on the “sweet spot” of Channel 2. Top: binned aperture flux measurements of three targets, normalized to the median of each target. Well depth was varied using instrument engineering requests (IERs) to vary the frame time and corrected with special dark frames obtained the same way. Data were binned by frame time. A 4th order polynomial fit to the data is overlaid as a dashed curve. Names of the stars are labeled, along with average flux densities. Bottom: Unbinned aperture flux vs. fluence data for 15 additional stars measured using AORs, color-coded by star. Frame times are labeled (in seconds) at the location of the median flux and fluence for a given frame time and star, with the qualifier “S” or “F” to denote full or sub-array readout mode. For comparison, we have overlaid the fit from the top plot. Measurements with perfectly corrected nonlinearity should lie on a straight horizontal line.
Figure 8.15: Peak pixel value in DN vs. distance from the sweet spot in channel 2. Sweet spot is defined here to be [15.12, 15.1] (when [0,0] is defined as the bottom left corner). These stars are all observed with multiple snapshots, and include the gain map calibration star itself taken with two different frame times. The second plot shows the different slope when using a 28 K source.
Figure 8.16: Bias Level in DN as a function of time. Bias level changes on long time scales (≈ 10 hours). Bias level as a function of time appears to not be repeatable from AOR to AOR. The data set used to make these plots consists of ten serendipitous staring mode AORs (five in each channel, 0.4 second frame time) in the off-field position, so theoretically empty of stars, although there are some stars in these observations. This bias change is not dependent on the frame delay (all frame delays show the change as a function of time). From these five AORs per channel, the maximum deviation is ≈ 2 DN in channel 1 and ≈ 0.5 DN in channel 2. These are binned over the whole BCD, so each point includes the mean from all 64 subframes. Only one amplifier is shown. Blue and black points show the different frame delays, for which there is no correlation with the shape of the long time scale change. Here we show two plots for each channel to demonstrate the range of level changes.
Figure 8.17: Channel 1 and 2 ensembles of primary calibrator fluxes as a function of time since the start of the mission. Black points show the ensemble of primary calibrators while green stars show the cosmic ray counts.
Figure 8.18: Individual calibration stars as a function of time. There is no evidence for a single outlier star as the source of the flux degradation trend.
Figure 8.19: A Monte Carlo simulation of calibration flux vs. time stamp slopes. This figure shows what the slopes would be if the time stamps for each data point were shuffled using a Fisher - Yates shuffle. The actual measured slope for channel 2 is shown with the dashed line. Because this measured slope is outside of 1σ, we take the slope to be significant.
Figure 8.20: Skydark variations in channel 1 over the course of the warm mission. The y-axis is a difference of the flux density present in the darks after an average dark frame has been removed.
Figure 8.21: Eclipse depths for ten visits to XO-3b, as computed via various methods. The group of points for each epoch is separated to minimize confusion. Error bars in this plot are symmetric; in cases where the technique returned asymmetric uncertainties, we used the largest of the two values. We show the results for the separate visits to the left of the gray vertical line, and the average results to the right. Error bars on the separate visits are the uncertainties reported by the technique. Error bars on the averages are the uncertainties in the weighted mean, adjusted for “underdispersion” using the scatter in the group of measurements. The horizontal red lines display the grand mean for all the results, and its uncertainty.
Figure 8.22: Similar to Figure 8.21. A blue horizontal line indicates the eclipse depth input to the simulations, 1875 ppm.
Figure 8.23: Actual precision reached for published, single epoch, IRAC observations. Plots are normalized at a binning size of one. The dashed black line is photon noise limited. The dashed grey line is two (1.5) times the photon noise limit for channel 1 (2). The third plot is a zoom-in on channel 2.
Appendix I. Appendix List of Figures
Figure C.1: PRF fits vs. aperture photometry for selected IRAC calibration star CBCDs. The vertical axis is the fractional difference between the PRF fit and corrected aperture photometry. The aperture photometry for IRAC channels 3 and 4 is in a 3-pixel radius with a 12 – 20 pixel background annulus and an aperture correction factor from this Handbook. For IRAC channels 1 and 2, it is in a 10-pixel radius with the same annulus. Short black dashed lines are the expected annulus correction needed. Long blue dashed line is the offset estimated from a weighted average of the data. Note this is essentially the expected value for IRAC channels 3 and 4. But IRAC channel 1 (and IRAC channel 2 to a lesser extent) requires a pixel phase correction (see text).
Figure C.2: Data from Fig. C.1, with IRAC channels 3 and 4 corrected for the annulus contribution, and IRAC channels 1 and 2 corrected for the pixel phase effect.
Figure C.3: APEX PRF-fitted photometry in the Serpens test field, with array location-dependent PRFs vs. aperture photometry. The aperture has a 3 pixel radius, the background annulus is 12 – 20 pixels. The aperture fluxes have been corrected using the aperture corrections in Table 4.8. The IRAC channel 3 and 4 PRF fluxes have been corrected for annulus contribution.
Figure C.4: APEX PRF-fitted photometry with a PRF Map vs. aperture photometry in the Serpens test field. PRF and aperture fluxes have been corrected as described in the text.
Figure C.5: The 25 PRF positions on an IRAC BCD.
Figure D.1: A BADTRIG problem in staring mode data. Left side vertical axis is for the grey lines and wheel speed RPM. The right-side vertical axis corresponds to x (blue) and y (green) centroids of the AOR as a function of sclk time. BADTRIG is set for the x-positions which have centroids of 14.5 after sclk time of 1.058017x109.
Figure D.2: Background measured in electrons as a function of time for 0.02 second subarray data.
Figure D.3: Same as Figure D.2, only now with each data point connected by a line, and zoomed in on the background modulation portion of the light curve.
Figure D.4: Same as Figure 4.13, except this AOR was during the 2013 anomaly. A wheel speed shift occurs at time 1.057755, but the centroids never return to their pre-wheel-shift positions.
Figure D.5: x- and y-centroids as a function of time. The pointing history file changes just beyond sclk time 1.06246x109, which changes the expected position of the star in the image, and therefore makes the centroiding algorithm essentially fail to find the target. This AOR has been re-processed and now no longer shows this problem.
Figure D.6: x- and y-positions as a function of time for a WASP-15 AOR. The position points which do not follow the sawtooth pattern are caused by CRs.
Figure D.7: Strong horizontal dark stripes in a 6-second frame time raw image in channel 2.