Appendix E. List of Figures
Figure 2.1: The IRS Cold Assembly.
Figure 2.2 Schematic representation of the IRS slits and Peak-Up apertures. Note that the IRS slits are not parallel in the Spitzer focal plane.
Figure 2.3: The PSF at 8 microns, centered in the SL slit (shown as a rectangle 3.6 arcsec wide). Simulated image using STINYTIM.
Figure 2.4: Short-Low spectral (grey) & Peak-Up (black) optical components and light paths.
Figure 2.5: Long-Low module optical components and paths.
Figure 2.6: Short-High module optical components and paths.
Figure 2.7: Long-High module optical components and paths.
Figure 2.8: PUI response functions, in electrons/photon. This is the product of the detector’s responsive quantum efficiency (QE times photoconductive gain) and the filter transmission.
Figure 2.9: PUI PSFs for the blue (top) and red (bottom) apertures. These PSFs are co-additions of several high signal-to-noise stellar images placed within 0.1 pixels of one another on the array. The images on the left are stretched to show the central PSF and the first diffraction ring; the images on the right are stretched to show the second diffraction ring. These images show the full fields of view of the peak-up arrays.
Figure 2.10: DQE for the Si:Sb arrays as a function of wavelength for different bias settings (1-2.5 V). The DQE declines linearly from 30 microns to 40 microns (reaching zero at approximately 40 microns) for the optimal bias voltage (1.6 V) that provides the best performance over the 14-36 micron interval.
Figure 2.11: Si detector output signal path (one of four parallel paths).
Figure 2.12: Double Correlated Sampling (DCS) data collection technique.
Figure 2.13: “Sample-Up-the-Ramp” data collection technique.
Figure 2.14: 1-sigma continuum sensitivity (in the faint source limit) of all four IRS modules in a 512 second integration with HIGH (top) and MEDIUM (bottom) backgrounds (the LOW background case, corresponding to an ecliptic latitude of 90o, is indistinguishable from the MEDIUM background case at the scale of these plots).
Figure 2.15: Short-Low 1-sigma point source staring continuum sensitivity (PSSCS) in the faint source limit after smoothing to a resolution of R = 50, with HIGH (top), MEDIUM (middle), and LOW (bottom) backgrounds.
Figure 2.16: Long-Low 1-sigma point source staring continuum sensitivity (PSSCS) in the faint source limit after smoothing to a resolution of R = 50, with HIGH (top), MEDIUM (middle), and LOW (bottom) backgrounds. This figure includes the effects of the LL1 filter delamination discovered before launch.
Figure 2.17: Short-High 1-sigma point source sensitivity (PSSL) to an unresolved emission line in the faint source limit, with HIGH (top) and MEDIUM (bottom) backgrounds (the LOW background case, corresponding to an ecliptic latitude of 90o, is indistinguishable from the MEDIUM background case at the scale of these plots).
Figure 2.18: Long-High 1-sigma point source sensitivity (PSSL) to an unresolved emission line in the faint source limit with HIGH (top) and MEDIUM (bottom) backgrounds (the LOW background case, corresponding to an ecliptic latitude of 90o, is indistinguishable from the MEDIUM background case at the scale of these plots).
Figure 2.19: Saturation levels for the shortest ramp duration available in the IRS AOTs, with MEDIUM background (the HIGH and LOW background cases, corresponding to ecliptic latitudes of 0o and 90o, respectively, are indistinguishable from the MEDIUM background case at the scale of this plot). For flux densities at a given wavelength that exceed these limits, at least one of the samples will be saturated. Flux densities up to 3 times greater will have at least two unsaturated samples.
Figure 2.20: Bright Source Limit (BSL) as a function of integration time for each of the IRS modules. Targets with fluxes above the curves for a given integration time are source-dominated; targets with fluxes below the curves are read noise or background limited. The plotted points on the curve for each module indicate the allowed AOT durations for that module.
Figure 4.1: An example of the co-added pre-flatfielded data for a stepped standard star in LL1 (left panel) and LL2 & LL3 (right panel) modules.
Figure 4.2: LL1 and LL2 co-added stellar data combined.
Figure 4.3 The model of a stellar spectrum in 2D created by interpolating the 1D spectrum over the wavsamp file.
Figure 4.4 The LL protoflat.
Figure 4.5 A cmask for the LL-PRE25 data. The pixels marked white are the ones labeled uncertain by our process.
Figure 4.6 The spatial correction for the LL example, generated by taking a running 5-row average of the protoflat / zodiacal background file. Multiplying this back into the flat gives us the final flat.
Figure 4.7 Left: The final flat for LL. Right: The zodiacal background divided by the final flat, as a check to ensure a spatially flat image.
Figure 4.8 Measured LL resolution. The scatter is due to PSF undersampling and pointing errors.
Figure 4.9 Measured SL resolution.
Figure 4.10: Comparison between standard and constant-width extractions for Short Low 2, nod 1.
Figure 4.11: Comparison between standard and constant-width extractions for Short Low 2, nod 2.
Figure 4.12 Comparison between standard and constant-width extractions for Short Low 1, nod 1.
Figure 4.13 Comparison between standard and constant-width extractions for Short Low 1, nod 2.
Figure 4.14: Comparison between standard and constant-width extractions for Long Low 2, nod 1.
Figure 4.15: Comparison between standard and constant-width extractions for Long Low 2, nod 2.
Figure 4.16: Comparison between standard and constant-width extractions for Long Low 1, nod 1.
Figure 4.17: Comparison between standard and constant-width extractions for Long Low 1, nod 2.
Figure 4.18 Left: A 2D image of the 16 micron PSF, based on observations of the star HD 42525. The relevant data have AORKEYs 17080064 and 17082112, and were mosaicked using MOPEX. The image is approximately 54 arcsec on a side. Right: A radial profile of the Blue Peak-Up PSF.
Figure 4.19 Left: A 2D image of the 22 micron PSF, based on observations of the star HD 172728. The relevant data have AORKEYs 20552960 and 20678912, and were mosaicked using MOPEX. The image is approximately 54 arcsec on a side. Right: A radial profile of the Red Peak-Up PSF.
Figure 4.20 Photometric accuracy of the red (top) and blue (bottom) peak-up imaging. PUI fluxes are compared to the fluxes derived from spectra of each star, integrated over the PUI bandpasses. The red PUI calibrators are HD163466 (23.4 mJy), HD172728 (36.9 mJy), HR5467 (37.3 mJy), HD173511 (254 mJy), and HR6348 (240 mJy). The blue PUI calibrators are HD163466 (46.6 mJy), HD172728 (73.7 mJy), and HR5467 (74.3 mJy).
Figure 5.1 IRS Pipelines Flow chart. The BCD Pipeline is described in Sections 5.1 and 5.2, the Peak-Up Cutouts Pipeline is described in Section 5.4, the Two-Dimensional Coadder Pipeline is described in Section 5.5, the Background Subtraction Pipeline is described in Section 5.6, and the Spectral Extraction Pipeline is described in Section 5.7.
Figure 5.2 BCD Pipeline Flowchart. The red line indicates a pathway the archive takes, but that is not offered in CUPID. The calibration files shown in dark are produced manually offline. See Section 5.3 for a description of the creation of dark.fits.
Figure 5.3 Before computing Poisson noise, the module SNEST_IRS applies an offset to each input ramp top panel. First, pair-wise differences for each sample are computed. Then the median of these (shown in the bottom panel) is used to exclude outlier samples, namely those whose . In this case, the first sample is an outlier. Finally, the ramp is offset (bottom panel such that the minimum non-outlier sample (in this case the second sample) has the value . This calculation is internal to SNEST_IRS and is not propagated. Outlier samples get assigned an arbitrarily large Poisson noise value.
Figure 5.4 This ramp exhibits A/D saturation in samples 8-16 (x). SATCOR outputs saturation-corrected samples (asterisks). The operation consists of first obtaining the mean slope of the unsaturated portion of the ramp (dashed line). This slope is then used to extrapolate the signal from the last unsaturated sample (indexed as np-1), to the saturated sample immediately following (np). The calculation takes into account non-linearity of a realistic ramp, so that the extrapolated sample np will be smaller than if it were along the dashed line. The extrapolation is done in succession for higher samples up to the top of the ramp. The saturation-corrected samples are used only by the subsequent modules DROOP and ROWDROOP.
Figure 5.5 Illustration of non-linearity correction in a LL 6 second ramp for a bright source: the open rhombii are input samples, referred to as Sobs in the text; the filled circles are output samples, referred to as Slin, after the operation of LINEARIZ. The solid line is the linear fit computed by the subsequent module SLOPE_FINDER.
Figure 5.6 Illustration of darkdrift correction in the top plane of a SL 14 second DCE. The left frame is the input lineariz.fits, while the right frame is the output darkdrift.fits. Note that the prominent jail bars in the input practically disappear in the output. The correction is done on each plane of the DCE separately.
Figure 5.7 Slope in the presence of a radhit: the radhit shown occurs in sample # 8 of the ramp. SLOPE_FINDER excludes it and subsequent (top) samples, because the slope after a radhit is shallower than the slope extrapolated (dashed line) from the bottom of the ramp. The output slope (solid line) is computed from the bottom samples preceding the radhit. Sample # 1 is always excluded, and the slope between samples # 4 and 5 in this ramp is excluded as an outlier, during the outlier rejection iteration. Shown also are slopes ±1 from the mean slope, where is the output uncertainty for this ramp.
Figure 5.8 Correcting the stray light. The frame on the left displays a SL ground-based laboratory image. The peak-ups are saturated and have been blacked out for maximum contrast. The white lines drawn on the image delineate the areas in the stray light mask used to determine the parameters bij. The frame on the right shows the same image, corrected for stray light.
Figure 5.9 Dark Calibration Pipeline Flow Chart. This pipeline consists of two steps: Pre-processing, up to DARKBASE, and ensemble processing, which combines several DARKBASE outputs (drk.fits). The figure above illustrates the case of three drk.fits that are combined to produce the output dark.fits.
Figure 5.10 Extraction Pipeline flowchart. The Extraction Pipeline runs after the BCD Pipeline and extracts one-dimensional spectra from two-dimensional slope images. The input of the thread is a two-dimensional BCD image in FITS format. The output from the thread is a table file containing the extracted spectrum. The thread also takes mask and uncertainty files as input, and propagates uncertainties and status flags into the output table.
Figure 5.11 One spectral order defined by the wavsamp.tbl file is illustrated, with a single pseudo-rectangle highlighted. Order curvature is highly exaggerated here for illustration.
Figure 5.12 An example of the plotted output of the PROFILE module.
Figure 5.13 (Left) A pseudo-rectangle is divided into pseudo-pixels. The pixels form the FITS image for a rectangular grid, as illustrated in the upper left corner. (Right) A pseudopixel is shown overlaid on the pixels of the original FITS image. The pseudo pixel contains partial detector pixels, so each partial pixel must be carefully accounted for to perform the extraction.
Figure 6.1: Schematic relationship between BCD files.
Figure 6.2 Schematic relationship among PBCD files.
Figure 6.3: Flat field images for the Short-Low (left) and Long-Low (right) modules. These images illustrate the full areas of the detector arrays dedicated to each spectral order. For both modules, the 1st order (longer wavelength) spectrum is on the left and the 2nd order (shorter wavelength) spectrum is on the right; wavelength increases from top to bottom within each order. The short bonus segment is visible above the 2nd order spectrum for both modules. The SL image also shows the two peak-up arrays. Spectral fringing is visible in the LL1 spectrum.
Figure 6.4: Unprocessed point-source data from the SL (left) and LL (right) modules. The SL image shows a spectrum obtained in 2nd order (along with the 1st order bonus segment at the top and the saturated IRS peak-up arrays on the right), while the LL image shows a spectrum obtained in 1st order.
Figure 6.5: Example bcd.fits image of a galactic nucleus. In this case the galaxy was observed using the 1st order Short-Low sub-slit (SL1). The bright continuum can be seen running vertically down the image on the left. Two bright (and several faint) emission lines/bands are clearly detected as features of increased intensity superimposed on the continuum. The square regions to the right are the well-exposed peak-up apertures. Note that light from the 2nd order SL sub-slit (SL2) is also deposited on the detector array to the right of SL1; however, SL2 is off-target (i.e., SL2 is looking at blank sky). This is because the SL2 sub-slit is located more than 1 arcminute away from the SL1 target position on the Spitzer focal plane. Consequently, there is no apparent SL2 spectrum.
Figure 6.6: Example bcd.fits and coa2d.fits IRS data products for a bright (~0.5 Jy at 15 microns) galaxy with strong emission lines observed using the Short-High (SH) module of the IRS. The rest frame 12.81 micron [NeII] line, appearing at an observed wavelength of ~13.1 microns, is circled in both images. Due to the overlap between orders in SH, the [NeII] emission line appears in both orders 15 and 16 in these data. The bcd.fits and coa2d.fits files are from the first nod position. Six individual nod frames were combined to produce the coa2d.fits image.
Figure 6.7: Unprocessed point-source data from the SH and LH modules. Both images show the typical layout of a cross-dispersed echellogram.
Figure 6.8 Example image of a star in the red IRS peak-up array. This image shows the onboard-processed DCE #3, with the rest of the image masked.
Figure 6.9: In-orbit image of a point source on the blue IRS peak-up sub-array. The first diffraction ring of the point source is clearly visible. Bright pixels in the image are cosmic ray hits. The illuminated red peak-up sub-array is visible in the figure above the blue sub-array, but does not contain any stars.
Figure 7.1: Latent charge build-up in electron/s on the LL array in a long (120 s) integration.
Figure 7.2: 129 observations of the standard star HD 173511, processed using the S18.18 pipeline. The observations used high accuracy self-peakup. In each case the spectra have been normalized to the median of all fluxes at nod 1.
Figure 7.3 Order mismatch between SL1 and LL2, which might be mistaken as a broad absorption feature. For each order the two nods are shown.
Figure 7.4 IRS Staring observations of 29 Vul, processed with S15 (bksub products). The red traces correspond to IRS observation campaign 7, while the blue correspond to five other campaigns.
Figure 7.5 Illustration of the 24 micron deficit. The 14 micron teardrop is also visible.
Figure 7.6 Another illustration of the 24 micron deficit. The 14 micron teardrop is also visible.
Figure 7.7 Example *bcd.fits image illustrating the 14 micron teardrop. The source was observed using SL1. The bright continuum can be seen running vertically down the image on the left. The white circle indicates the teardrop, which appears as a region of excess flux to the left of the continuum.
Figure 7.8 Illustration of 14 micron teardrop in the extracted spectra of HR 7341. The top two panels show the results for Regular Point Source extraction and the bottom two panels show the results for Full Width extraction. The teardrop is most clearly seen in panels 2 and 4, where it appears as an upturn kicking in at around 13.5 microns.
Figure 7.9 Same as Figure 7.8, but for HR 2194.
Figure 7.10 Similar to Figure 7.8, but for 3C 454.3, a red source. Comparisons to models have been omitted due to unavailability of models.
Figure 7.11 LNZ Cube: Output of linearize – Frames 1-4.
Figure 7.12 Various pipeline products which all show the 14 micron teardrop: After the various droop corrections (top left); after stray light correction but before flat fielding (top right); after flat fielding (lower left); flat fielded but NOT stray light corrected (lower right).
Figure 7.13 Nine mapping positions (2, 12, 22, 27, 32, 37, 42, 52, and 62) of the flat fielding data for eta1 Dor, zoomed in on the teardrop. The teardrop shape changes slightly with slit position, consistent with an internal reflection origin.
Figure 7.14 Two dimensional SL2 spectral image showing an observation of HR 7341. The wavelength region spanning 6.6-7.5 microns is highlighted with a box. We do not see any excess emission similar to the teardrop between 6.6 and 7.5 microns in SL2. This suggests the feature in SL2 at 6.5-7.6 microns is different in nature from the SL1 teardrop, and likely does not depend on extraction aperture.
Figure 7.15 The strength of the teardop (quantified as the ratio of flux densities at 14.0 and 13.5 microns) versus the source flux density (measured at 12 microns) for 8 sources. For each source, black symbols represent full-slit extractions and red symbols represent point-source extractions. Plus signs represent nod 1 and asterisks represent nod 2. There is no correlation between the strength of the teardrop and the flux of the source. However, the teardrop is almost always stronger at nod 1 than nod 1 for point source extraction.
Figure 7.16 Illustration of excess LH dark current, which is distributed unevenly over the array with the brightest region towards the blue end (bottom) of the array.
Figure 7.17 Nod 1 and nod 2 spectra. Notice the blue tilts in the nod 1 spectrum.
Figure 7.18 The LH spectrum of Mrk 231 shows red tilts in orders 1-3 (29-37 micron).
Figure 7.19 A bcd from AOR 16921344 of the planetary nebula J900. Red circles highlight the locations of the spectral ghost features.
Figure 7.20 LH spectrum. Notice the peaks at 19, 20, and 21 microns in the S13.2 spectrum.
Figure 8.1 This flowchart describes the basic steps necessary to reduce IRS low-resolution spectra. Additional flowcharts illustrating the 2D background subtraction procedure and how to clean up the 1D spectrum are provided in Figure 8.2 and Figure 8.3.
Figure 8.2 This flowchart illustrates how to subtract the background from your low-resolution 2D IRS spectrum.
Figure 8.3 This flowchart illustrates the main issues found with low-resolution IRS spectra. These include latent charge (see Section 7.2.5), residual bad pixels (which can be corrected using IRSCLEAN; see Section 7.2.2), and fringes (which can be corrected using IRSFRINGE or SMART; see Section 7.3.7).
Figure 9.1 - Example IRS Merged Spectrum (black line). The IRAC 8 (blue), IRAS 12 (cyan), IRS PU 16 (green), IRS PU 22 (orange), MIPS 24 (red), and IRAS 25 (blue-green) transmission curves are superimposed on the merged spectrum. The synthetic flux in each band is indicated by a colored horizontal line. For clarity, the MIPS 24 and IRAS 25 photometry have been shifted up or down by 0.002 Jy, respectively.
Figure 9.2 - Comparison of median synthetic photometry of IRS flux standard stars with actual IRAC 8, MIPS 24, PU 16, and PU 22 measurements of these same stars. These stars were observed many times by Spitzer IRS. The vertical error bars represent the standard deviation of the mean over all visits for the IRS spectrophotometry. The slopes of the best-fit (dotted) lines give an indication of how well the synthetic photometry compares to the integrated light measurements of the same stars. Unity slopes are indicated by the solid lines.
Figure 9.3 - MIPS 24 versus IRAC 8 synthetic IRS spectrophotometry. A narrow locus of stars with blue (Rayleigh-Jeans) MIPS 24/IRAC 8 colors of ~0.2 and a broader band of galaxies with red colors in the range of 1-10 are both apparent. IRS calibration stars (blue diamonds) and galaxies and AGNs (red diamonds) are shown for reference.
Figure 9.4 - Distribution of IRS position measurements for IRS calibration sources, relative to requested position. Circles with radii of 0.36”, 0.62”, and 1.35” enclose 68%, 95%, and 99% of the measurements, respectively. The three X’s mark observations taken during Campaign 7, when observations were affected by known telescope pointing issues. (Nine additional deviant observations from this campaign fall off of the plot.)
Figure 9.5 - Distributions of IRS spatial profile widths for the SL2, SL1, LL2, and LL1 orders. The profile widths generally increase going from short to long wavelengths, reflecting the spatial resolution of the telescope and instrument. Distributions are shown for blue sources (blue lines) and red sources (red lines), separated at a MIPS 24/IRAC 8 flux density ratio ratio of 1 (Fig. 9.3). Calibration stars (not shown) follow a distribution similar to that of the blue sources. Galaxies tend to be red sources with MIPS/IRAC8>1.
Appendix F. List of Tables
Table 2.1: IRS module characteristics.
Table 2.2 Summary of changes in bias and temperature for all IRS modules.
Table 2.3 SUR data collection mode sampling parameters.
Table 2.4 PUI 1σ sensitivity (microJy) for 6/14/30 second ramp times.
Table 4.1 Data used to create IRS flats.
Table 4.2 Lines used for Wavelength Calibration.
Table 4.3: IRS Wavelength Calibration.
Table 4.4 AORs used to calibrate S18.18 SL data, PRE-33 campaigns.
Table 4.5 AORs used to calibrate S18.18 SL data, POST-32 campaigns.
Table 4.6 AORs used to calibrate S18.18 LL data, PRE-45 campaigns.
Table 4.7 AORs used to calibrate S18.18 LL data, POST-44 campaigns.
Table 4.8 AORs used to calibrate S18.18 SH data, all campaigns.
Table 4.9 AORs used to calibrate S18.18 LH data, PRE-25 campaigns.
Table 4.10 AORs used to calibrate S18.18 LH data, POST-24 campaigns.
Table 4.11: Aperture Corrections to Infinity (1.8 arcsec/pix).
Table 4.12 Peak-up Imaging Color Corrections for Blackbodies.
Table 4.13 Peak-Up Imaging Color Corrections for
Table 5.1: Sample wavsamp.tbl file.
Table 5.2: Sample profile.tbl file (excerpt).
Table 6.1 Classes of Spitzer data products.
Table 6.2 Number of reads (samples) per ramptime available for each module.
Table 6.3 Raw Files Delivered.
Table 6.4: Calibration Files Delivered.
Table 6.5: Wavelength ranges used to derive fluxcon.tbl.
Table 6.6 BCD Files Delivered.
Table 6.7 PBCD Files Delivered.
Table 6.8 Peak-Up Acquisition (PUA) Files Delivered.
Table 6.9: PUI Files Delivered.
Table 6.10: pmask Bit Settings.
Table 6.11: dmask Bit Settings.
Table 6.12: bmask Bit Settings.
Table 6.13: c2msk/bkmsk Bit Settings.
Table 6.14 Translation between header keyword values for FOVID and FOVNAME.
Table 9.1 – Orders and wavelengths in a Merged Spectrum file.
Table 9.2 - Data columns in a Merged Spectrum table file:
Table 9.3 - Distribution of spatial profile widths:
Table 9.4 - Orders and wavelengths in a Merged Spectrum file:
Table 9.5 - Columns in the Catalog :