Creation of Model Particle Background Spectra and Images

The creation of the quiescent particle background (QPB) spectra is a two step process requiring mosspectra and mosback or pnspectra and pnback. The first step can be rather time consuming (though we have implemented a new way of making the process faster) while the second step runs quite quickly.

The mosspectra and pnspectra tasks start the process of implementing the Kuntz & Snowden (2008) particle background calibration. They extract QPB spectral information from three sources: the corner spectra from each chip of the observation being analyzed, the corner spectra from each chip of the FWC data, and the FOV spectra from each chip of the FWC data. The FWC data are derived from observations where the filter wheel is in the closed position, blocking the detection of any X-rays from outside the instrument and the penetration of any soft protons. Cosmic rays, however, still penetrate to the detectors so a clean measure of the QPB is allowed. The FWC data provide the spatial distribution of the QPB events. The corner data provide a measure of the both normalization of the QPB rate and the shape of the QPB spectrum, which was found by Kuntz & Snowden (2008) to be temporally variable.

While mosspectra and pnspectra extract these data, mosback and pnback combine all of these different pieces into the QPB spectrum for a specified region in the FOV for the observation and an image of the QPB component in a given band. Because the count rate in the corner data is quite low, the corner spectrum is augmented from a database of corner data extracted from the public archive. (See Kuntz & Snowden, 2008 for a full explanation.) This augmentation is done by mosback and pnback.

Figure 12: Upper Panel: Observed spectrum (red) and model QPB spectrum (green) for the full field of view for the MOS1 observation of the Abell 1795 cluster observation from ObsID 0097820101. Note the discrepancy at high energies where the model QPB is still low compared to the observed spectrum even though there is essentially no effective area for detecting cosmic X-rays. This is another indication of residual SP contamination as first suggested by the low-level variations in the light curve for the nominally low intensities. Note the emission lines including the strong Fe line at $\sim6.5$ keV in the observed spectrum of Abel 1795. The smooth part of the model QPB at $\sim1.5$ keV is the “bridge” where the Al K$\alpha$ and Si K$\alpha$ fluorescent instrumental lines affect the data. Lower Panel: The complete FWC spectrum from the MOS1 instrument showing the very strong Al K$\alpha$ and Si K$\alpha$ fluorescent instrumental lines ($\sim1.49$ keV and $\sim1.75$ keV, respectively) on top of the continuum. Other fluorescent lines are visible at higher energies as well as a strong low-energy tail due to detector noise.

Spectra: Figure 12 (upper panel) displays the observed spectrum and model QPB spectrum from the MOS1 observation of Abell 1795 from the full field of view. The observed spectrum clearly dominates over the background over the entire energy range where the MOS has significant response (i.e., $<8.0$ keV). In spite of the fact that the Abell 1795 cluster is relatively hot and bright, the excess of the observed spectrum over the model QPB spectrum at energies above $\sim10$ keV is not cosmic in origin. The excess is most likely due to residual SP contamination, as suggested by the low-level variation in the accepted time intervals of the light curve (the green part of the curves in Figure 8).

Figure 12 (lower panel) displays the FWC spectrum from the MOS1 instrument due to the QPB. The most striking features are the low-energy tail, the strong Al K$\alpha$ and Si K$\alpha$ fluorescent lines near 1.49 keV and 1.75 keV, and several other fluorescent lines at higher energies. The Al and Si lines are the only problematic parts of the spectrum above 0.3 keV. Because these lines are so strong, any errors or slight changes in the intrinsic line spread function will show up as significant residuals. Since the FWC data is collected over the entire mission, and the line spread function has changed over that time, the FWC data do not provide a usable template for the background at the Al K$\alpha$ and Si K$\alpha$ energies. Small gain variations can also produce very striking residuals in spectral fits where the surface brightness of the object of interest is typically smaller than the strength of the fluorescent lines. This has not been observed to be a problem for the higher energy lines. To circumvent this problem, a smooth “bridge” is fit to the data on either side of the Al K$\alpha$ and Si K$\alpha$ lines and the continuum component interpolated. This “bridge” is the smooth part of the model QPB shown in Figure 12. As discussed below in § 6, the Al K$\alpha$ and Si K$\alpha$ lines must then modeled as separate Gaussians in spectral fits.

Below 0.3 keV, the QPB rises rapidly. Thus we do not recommend attempting spectral fitting below 0.3 keV.

Figure 13: Upper Panel: Observed spectrum (red) and model QPB spectrum (green) for the full field of view for the pn observation of the Abell 1795 cluster. Note that besides the strong Fe emission line at $\sim6.5$ keV there are strong instrumental Cu, Cr, and Ni lines in the observed spectrum. The smooth parts of the model QPB at $\sim1.5$ keV and 8 keV are the bridges where the Al K$\alpha$ and higher energy (note that the pn is not affected by a Si K$\alpha$ line) fluorescent instrumental lines strongly affect the data. Lower Panel: The complete filter wheel closed spectrum from the pn instrument showing the very strong Al K$\alpha$ and Cu, Cr, and Ni fluorescent instrumental lines ($\sim1.49$ keV and $\sim8$ keV, respectively) on top of the continuum. Other fluorescent lines are visible at higher energies as well as a strong low-energy tail due to detector noise.

Figure 13 (upper panel) displays the observed spectrum and model QPB spectrum from the pn observation of the cluster Abell 1795. Similar to Figure 12, the cluster spectrum dominates at low energies. However, unlike the MOS spectrum, the pn spectrum has several very strong instrumental Cu, Cr, and Ni lines in the 7.2-9.0 keV band and is missing the Si K$\alpha$ line at 1.75 keV. With such strong Cu lines it may be reasonable to simply mask those data from the spectral analysis. Similar to MOS spectral fits, the Al K$\alpha$ and Cu, Cr, and Ni lines in the $7.2-10.0$ keV range must be explicitly fit. Doing so, spectral data for hard sources can be used up to $\sim11$ keV.

As noted previously, the lower energy limit depends upon the pattern selection. Selecting only PATTERN$==$0 allows one to reduce the QPB contribution at low energies and one can extract a spectrum down to 0.3 keV to reasonable confidence. Selecting PATTERN$==$0 has the consequence of significantly reducing the count rate at higher energies. Selecting PATTERN$<=$4 provides a higher count rate at higher energies, but produces a strong QPB signal at low energies. With PATTERN$<=$4 it would be unwise to extract the spectrum below $\sim$0.55 keV.

Figure: Output MOS1 images from mosspectra of the observed counts (upper panel) and model exposure (lower panel) in sky coordinates. The images are for the $0.35-1.25$ keV band.

Imaging: Figure 14 shows typical output images for a MOS exposure: the total counts in the FOV (P-fovimsky-elo-ehi.fits, made by mosspectra), the exposure map (P-expimsky-elo-ehi.fits, made by mosspectra), the quiescent particle background image (P-bkgimdet-elo-ehi.fits, made by mosback) and the source mask (P-cheeset.fits, made by cheese). The residual soft proton image, made by proton, and the solar wind charge exchange image, made by swcx, will be shown in §7.