Chapter 6
ACIS: Advanced CCD Imaging Spectrometer

6.1 Introduction & Layout

The Advanced CCD Imaging Spectrometer (ACIS) offers the capability to simultaneously acquire high-resolution images and moderate resolution spectra. The instrument can also be used in conjunction with the High Energy Transmission Grating (HETG) or Low Energy Transmission Grating (LETG) to obtain higher resolution spectra (see Chapters 8 and 9). ACIS contains 10 planar, 1024 x 1024 pixel CCDs (Figure 6.1); four arranged in a 2x2 array (ACIS-I) used for imaging, and six arranged in a 1x6 array (ACIS-S) used either for imaging or as a grating readout. Currently any combination of up to 6 CCDs may be operated simultaneously. However, see the discussion in Section 6.20.1 where we encourage users, where possible, to specify as many as 5 chips as optional. Prior popular combinations include an extended ACIS-I imaging mode, using chips I0-I3 plus S2 and S3; and an ACIS-S imaging mode, using chips S1-S4 plus I2 and I3. Operating six chips enhances the chance of serendipitous science, but at the price of increased total background counting rate and therefore a somewhat enhanced probability of saturating telemetry. Another penalty for 6 operating CCDs is an increased power load (Section 6.20.1). Two CCDs are back-illuminated (BI) and eight are front-illuminated (FI). The response of the BI devices extends to energies below that accessible to the FI chips. The chip-average energy resolution of the BI devices is better than that of the FI devices.

Figure 6.1: A schematic drawing of the ACIS focal plane; insight to the terminology is given in the lower left. Note the nominal aimpoints: on S3 (the `+') and on I3 (the `x'). It is standard practice to add an offset to all observations on S3 to move the source away from the node 0-1 boundary (see Section 6.10). Note the differences in the orientation of the I and S chips, important when using subarrays (Section 6.12.1). Note also the (Y, Z) coordinate system and the target offset convention (see Chapter 3) as well as the SIM motion (+/-Z). The view is along the optical axis, from the source toward the detectors, (-X). The numerous ways to refer to a particular CCD are indicated: chip letter+number, chip serial number, and ACIS chip number. The node numbering scheme is illustrated lower center.

Link to postscript file for Figure 6.1

Figure 6.2: A schematic drawing of the ACIS focal plane, not to scale. The ACIS-I array consists of chips I0-I3. Chips S2 and S3 are often turned on in addition to I0-I3 when ACIS-I is used in imaging mode as indicated in the top figure. ACIS-S consists of chips S0-S5. When operated in imaging mode chips S1-S4 plus I2 and I3 are often turned on (bottom figure). When operated in spectroscopy mode S0-S5 are usually turned on. See 6.20.1 for more details on chip choices.

Link to postscript file for Figure 6.2

The original Instrument Principal Investigator for ACIS is Prof. Gordon Garmire (Pennsylvania State University). ACIS was developed by a collaboration between Penn State, the MIT Kavli Institute for Astrophysics and Space Research and the Jet Propulsion Laboratory, and was built by Lockheed Martin and MIT. The MIT effort was led by Dr. George Ricker. The CCDs were developed by MIT's Lincoln Laboratory.

ACIS is a complex instrument having many different characteristics and operating modes. Radiation damage suffered by the FI chips has had a negative impact on their energy resolution - the BI devices were not impacted - thus impacting the basic considerations as to how to make best use of the instrument. We discuss the trade-offs in this chapter. Software methods for improving the energy resolution of the FI CCDs are discussed in Section 6.7.1. The low energy response of ACIS has also been affected by the build-up of a contaminant on the optical blocking filters and this is discussed in Section 6.5.1.

Many of the characteristics of the ACIS instrument are summarized in Table 6.1.

Focal plane arrays:
I-array	4 CCDs placed to lie tangent to the focal surface
S-array	6 CCDs in a linear array tangent to the grating Rowland circle
CCD format	1024 by 1024 pixels
Pixel size	23.985 microns (0.4920±0.0001 arcsec)
Array size	16.9 by 16.9 arcmin ACIS-I
	8.3 by 50.6 arcmin ACIS-S
On-axis effective Area	110 cm² @ 0.5 keV (FI)
(integrated over the PSF	600 cm² @ 1.5 keV (FI)
to > 99% encircled energy)	40 cm² @ 8.0 keV (FI)
Quantum efficiency	> 80% between 3.0 and 6.5 keV
(frontside illumination)	> 30% between 0.7 and 11.0 keV
Quantum efficiency	> 80% between 0.8 and 5.5 keV
(backside illumination)	> 30% between 0.4 and 10.0 keV
(QEs do not include contaminant
layer transmission)
Charge transfer inefficiency(parallel)	FI: ∼ 2×10⁻⁴; BI: ∼ 2×10⁻⁵
Charge transfer inefficiency(serial)	S3(BI): ∼ 7×10⁻⁵; S1(BI): ∼ 1.5×10⁻⁴;
	FI: < 2 ×10⁻⁵
System noise	< ∼ 2 electrons (rms) per pixel
Max readout-rate per channel	∼ 100 kpix/sec
Number of parallel signal channels	4 nodes per CCD
Pulse-height encoding	12 bits/pixel
Event threshold	FI: 38 ADU ( ∼ 150-350 eV)
	BI: 20 ADU ( ∼ 150-220 eV)
Split threshold	13 ADU
Max internal data-rate	6.4 Mbs (100 kbs ×4 ×16)
Output data-rate	24 kb per sec
Minimum row readout time	2.8 ms
Nominal frame time	3.2 sec (full frame)
Allowable frame times	0.2 to 10.0 s
Frame transfer time	40 μsec (per row)
Point-source sensitivity	4 ×10⁻¹⁵ergs cm⁻² s⁻¹ in 10⁴ s
	(0.4-6.0 keV)
Detector operating temperature	−90 to −120^°C

Table 6.1: ACIS Characteristics

6.2 Basic Principles

A CCD is a solid-state electronic device composed primarily of silicon. A "gate" structure on one surface defines the pixel boundaries by alternating voltages on three electrodes spanning a pixel. The silicon in the active (depletion) region (the region below the gates wherein most of the absorption takes place) has an applied electric field so that charge moves quickly to the gate surface. The gates allow confined charge to be passed down a "bucket brigade" (the buried channel) of pixels in parallel to a serial readout at one edge by appropriately varying ("clocking") the voltages in the gates.

The ACIS front-illuminated CCDs have the gate structures facing the incident X-ray beam. Two of the chips on the ACIS-S array (S1 and S3) have had treatments applied to the back sides of the chips, removing insensitive, undepleted bulk silicon material and leaving the photo-sensitive depletion region exposed. These are the BI chips and are deployed with the back side facing the HRMA.

Photoelectric absorption of an X-ray in silicon results in the liberation of a proportional number of electrons: an average of one electron-hole pair for each 3.7 eV of energy absorbed. Immediately after the photoelectric interaction, the charge is confined by electric fields to a small volume near the interaction site. Charge in a FI device can also be liberated below the depletion region, in an inactive substrate, from where it diffuses into the depletion region. This charge may easily appear in two or more pixels.

Good spectral resolution depends upon an accurate determination of the total charge deposited by a single photon. This in turn depends upon the fraction of charge collected, the fraction of charge lost in transfer from pixel to pixel during read-out, and the ability of the readout amplifiers to measure the charge. Spectral resolution also depends on read noise and the off-chip analog processing electronics. The ACIS CCDs have readout noise less than 2 electrons RMS. Total system noise for the 40 ACIS signal chains (4 nodes/CCD) ranges from 2 to 3 electrons (rms) and is dominated by the off-chip analog processing electronics.

The CCDs have an "active" or imaging section (see Figure 6.1) which is exposed to the incident radiation and a shielded "frame store" region. A typical mode of the ACIS CCD operation is: (1) the active region is exposed for a fixed amount of time (full frame ∼ 3.2 s); (2) at the end of the exposure, the charge in the active region is quickly ( ∼ 41 ms) transferred in parallel into the frame store; (3) the next exposure begins; (4) simultaneously, the data in the frame store region is transferred serially to a local processor which, after removing bias (see Section 6.13), identifies the position and amplitude of any "events" according to a number of criteria depending on the precise operating mode. These criteria always require a local maximum in the charge distribution above the event threshold (see Table 6.1). The position and the amount of charge collected, together with similar data for a limited region containing and surrounding the pixel are classified ("graded") and then passed into the telemetry stream.

6.3 Optical Blocking Filter & Optical Contamination

Since the CCDs are sensitive to optical as well as X-ray photons, optical blocking filters (OBFs) are placed just over the CCDs between the chips and the HRMA. The filters are composed of polyimide (a polycarbonate plastic) sandwiched between two thin layers of aluminum. The nominal thicknesses of the these components for the two arrays are given in Table 6.2. Details of the calibration of these filters may be found in the ACIS calibration report at http://www.astro.psu.edu/xray/docs/cal_report/node188.html. These calibrations do not include the more recent effects of molecular contamination. This is discussed in Section 6.5.1.

Table 6.2: Nominal Optical Blocking Filter Composition and Thicknesses

ACIS-I	Al/Polyimide/Al	1200Å 2000Å 400Å
ACIS-S	Al/Polyimide/Al	1000Å 2000Å 300Å

The threshold for optical contamination (a 1 ADU (3.4 eV) shift in the bias level) is based on on-orbit calibrations of a number of stars with different optical spectra. The threshold for detectable visible light contamination varies according to source color and is lowest for red stars observed on ACIS-S The detection threshold for an M star on the ACIS-S array is V ∼ 8.1 for the nominal 3.2 second exposure or V ∼ 6.3 using a 0.4s frame time and a 1/8 chip subarray. The thresholds are about 5 visual magnitudes brighter for the ACIS-I array. While the impact of the OBF molecular contamination has not been studied in detail, a preliminary estimate (using typical optical constants for organic materials) suggests little change in the broadband optical transmission. See Section 6.19 for a discussion of the observation of optically bright targets such as solar system objects.

6.4 Calibration

Measurements of ACIS include laboratory calibrations, a system-level ground calibration of the HRMA and ACIS at the X-Ray Calibration Facility (XRCF) at MSFC, and on-orbit calibration using celestial and on-board radioactive X-ray sources. The ACIS external calibration source (ECS) consists of an Fe-55 source and a target made of aluminum and titanium. The source emits five strong lines (Al K-α at 1.49 keV, Ti K-α and β, at 4.51 and 4.93 keV, and Mn K-α and β at 5.90 and 6.49 keV. A number of weaker lines are also present.

The on-orbit calibration of ACIS is a continuing activity. All calibration data are, or will be, described in detail, at http://cxc.harvard.edu/cal/. The user is urged to consult the WWW site and its links for the latest information.

6.5 Quantum Efficiency and Effective Area

The quantum efficiencies near the readout for the ACIS CCDs for the standard grade set, including optical blocking filters and molecular contamination, are shown in Figure 6.3. Note that the quantum efficiency for the FI chips varies somewhat with row number (not shown), and decreases by 5-15% farthest from the readout at energies above about 4 keV. This is due to the migration of good grades to bad grades produced by charge transfer inefficiency (CTI), which varies with row number. The quantum efficiency (QE) variation with position for the BI chips is much smaller.

Cosmic rays tend to cause large blooms on FI chips, and much smaller ones on BI chips. This results in a 2-4% decrement of the QE for FI chips and ∼ 0.5% for BI chips.

The combined HRMA/ACIS on-axis effective areas are shown in Figure 6.4 (log energy scale) and 6.5 (linear energy scale). The effective areas are for an on-axis point source and a 20-arcsecond-diameter detection cell. The ACIS effective areas include a correction that accounts for the build-up of molecular contamination on the ACIS filters (see the discussion in Section 6.5.1). Figures 6.4 and 6.5 show the predicted ACIS effective areas for the middle of Cycle 14 based on the current time-dependent ACIS contamination models.

There is a count rate dependent effect, currently under investigation, that may impact a subset of observers. The problem is that the effective area calibration at the Si K edge (1.842 keV; 6.74 Å) appears to be somewhat incorrect, but only at very high count rates. At high rates, the edge will appear somewhat higher than it should be. Here high count rate is: 1.9×10¹⁰ erg cm⁻² s⁻¹ (ACIS) and 1.1×10⁻⁹ erg cm⁻² s⁻¹ (HETG/ACIS). Of course at these counting rates pileup will likely dominate ACIS observations, but not necessarily HETG/ACIS observations. In very extreme cases, the effective discrepancy between the fluxes determined by the nominal effective area and observation may be as large as 10%. The effect manifests itself in strong residuals around the edge region for sources with moderate absorption columns ( < 3×10²¹ cm⁻²), an underestimate of the actual source column at higher absorption columns, and a change in effective area at higher energies.

Figure 6.6 shows the vignetting (defined as the ratio of off-axis to on-axis effective area) as a function of energy at several off-axis angles. These data are from a calibration observation of G21.5, a bright supernova remnant. The detector was appropriately offset for each off-axis angle so that the data were obtained at the same focal position, minimizing the effects of any spatially-dependent variations in the CCD response.

Figure 6.3: The quantum efficiency, convolved with the transmission of the appropriate optical blocking filter, of the FI CCDs (from a row nearest the readout) and the two BI CCDs as a function of energy. S3 is somewhat thicker, hence more efficient, than S1. These curves include the effects of molecular contamination, as discussed in the text.

Link to postscript file for Figure 6.3

Figure 6.4: The HRMA/ACIS predicted effective area versus the energy on a log scale. The dashed line is for the FI CCD I3, and the solid line is for the BI CCD S3. These curves include the effects of molecular contamination, as discussed in the text.

Link to postscript file for Figure 6.4

Figure 6.5: The HRMA/ACIS predicted effective area versus the energy on a linear scale. The dashed line is for the FI CCD I3, and the solid line is for the BI CCD S3. These curves include the effects of molecular contamination, as discussed in the text.

Link to postscript file for Figure 6.5

Figure 6.6: Vignetting (the ratio of off-axis to on-axis effective area) as a function of energy for several off axis angles in arcminutes.

Link to postscript file for Figure 6.6

6.5.1 Molecular Contamination of the OBFs

Astronomical observations and data acquired from the on-board ACIS external calibration source (ECS) show that there has been a slow and continuous degradation in the ACIS effective area at low energies. Our best explanation is that this is due to the build-up of out-gassed material condensing on the cold ACIS filters. The HRC operates at a warmer temperature and shows no sign of contamination build-up. The degradation in effective area is most severe at low energies. The build-up of contamination has been monitored by ECS observations before and after each perigee passage and by periodic observations with the gratings. The grating observations have shown that the contaminant is dominated by carbon, with smaller amounts of oxygen and fluorine. The ECS is used to monitor the degradation in effective area at energies above the C-K edge. The depth of the contaminant increases from the center of the detectors toward the edges where the filters are cooler.

An initial model of the contamination on the ACIS filter was developed for use in CIAO and other data reduction packages and released in CALDB v3.0.0 (Dec. 2004). Beginning in 2006, there were some systematic differences in the actual build-up rate of the contaminant on the ACIS filters compared with the predictions of the 2004 model. These differences were initially noticed from discrepancies in the optical depth of the contaminant as measured from ECS data and gratings observations of AGN. As a result, updated versions of the ACIS-I and ACIS-S contamination models were released in CALDB 4.2 (Dec. 2009). A further update to the ACIS-I contamination model was released in CALDB 4.4.1 (Dec. 2010). The new ACIS contamination models were derived from a re-analysis of all ECS data taken since launch. The ECS data measures the optical depth of the contaminant at 660 eV and Al-Ka (1.49 keV) as a function of time and position on the ACIS detectors. In addition, a large number of LETG/ACIS-S observations were also used in developing the models. These observations measure the optical depth of the contaminant at C-K, O-K and F-K edges. Finally, the new contamination models were extensively tested with observations of clusters of galaxies, supernova remnants and AGN. The new models produce more consistent, time-independent fluxes for clusters of galaxies and supernova remnants compared to the previous models. However, calibration observations taken in 2011 and a more extensive analysis of the external calibration source measurements suggest that the model of the contaminant's spatial pattern on the ACIS-I filter will have to be updated in a ∼ 200 pixel wide region along the border of the ACIS-I array. Users should avoid this region in the analysis of recent observations, if possible. The update will reduce the throughput in this region by several percent at 0.75 keV and have very little impact at 1.5 keV.

6.6 Spatial Resolution & Encircled Energy

The spatial resolution for on-axis imaging with HRMA/ACIS is limited by the physical size of the CCD pixels (24.0 μm square ∼ 0.492 arcsec) and not the HRMA. This limitation applies regardless of whether the aimpoint is selected to be the nominal point on I3 or S3 (Figure 6.1). Approximately 90% of the encircled energy lies within 4 pixels (2 arcsec) of the center pixel at 1.49 keV and within 5 pixels (2.5 arcsec) at 6.4 keV. Figure 6.7 shows an in-flight calibration. There is no evidence for any differences in data taken with either S3 or I3 at the nominal focus. The ACIS encircled energy as a function of off-axis angle is discussed in Chapter 4 (see Section 4.2.2 and Figure 4.13).

Figure 6.7: The on-orbit encircled broad-band energy versus radius for an ACIS observation of point source PG1634-706 (ObsID 1269). The curve is normalized to unity at a radius of 5 arcseconds. We estimate that the statistical uncertainty is 3%, and the systematic uncertainty due to power beyond 5^′′ is 2%. The effective energy is 1 keV.

Link to postscript file for Figure 6.7

Off-axis, the departure of the CCD layout from the ideal focal surface and the increase of the HRMA PSF with off-axis angle become dominating factors. Since the ideal focal surface depends on energy, observers for whom such considerations may be important are urged to make use of the MARX simulator to study the impact on their observation.

6.7 Energy Resolution

The ACIS FI CCDs originally approached the theoretical limit for the energy resolution at almost all energies, while the BI CCDs exhibited poorer resolution. The pre-launch energy resolution as a function of energy is shown in Figure 6.8. Subsequent to launch and orbital activation, the energy resolution of the FI CCDs has become a function of the row number, being near pre-launch values close to the frame store region and substantially degraded in the farthest row. An illustration of the dependence on row is shown in Figure 6.9.

Figure 6.8: The ACIS pre-launch energy resolution as a function of energy.

Link to postscript file for Figure 6.8

The loss of energy resolution is due to increased charge transfer inefficiency (CTI) caused by low energy protons encountered during radiation belt passages and Rutherford scattering through the X-ray telescope onto the focal plane. Subsequent to the discovery of the degradation, operational procedures were changed: ACIS is not left at the focal position during radiation belt passages. Since this procedure was initiated, no further degradation in performance has been encountered beyond that predicted from pre-launch models. The BI CCDs were not impacted, which is consistent with the proton-damage scenario - it is far more difficult for low-energy protons from the direction of the HRMA to deposit their energy in the buried channels of the BI devices, since the channels are near the gates and the gates face in the direction opposite to the HRMA. Thus the energy resolution for the two BI devices remains nearly at pre-launch values. The position-dependent energy resolution of the FI chips depends significantly on the ACIS operating temperature. Since activation, the ACIS operating temperature has been lowered in steps and is now set at the lowest temperature thought safely and consistently achievable ( ∼ −120^°C).

6.7.1 Correcting the Energy Resolution of the CCDs

The ACIS instrument team has developed a correction algorithm that recovers much of the lost energy resolution the FI CCDs. The correction recovers a significant fraction of the CTI-induced loss of spectral resolution of the FI CCDs at all energies. The algorithm has been incorporated in the CIAO tool acis_process_events as of CIAO 2.3. Figure 6.9 illustrates the improvement that the tool provides. As of December 2006, data for the two BI chips can also be corrected in the same way, including a correction for serial CTI for the BI chips, though the effects are more subtle. The resulting response is very nearly uniform across the BI chips once this correction is made. An alternative corrector may be found at the Chandra contributed SW exchange web site and also at the Penn State ACIS page (http://www.astro.psu.edu/users/townsley/cti/install.html).

Figure 6.9: We plot energy resolution (FWHM in eV) vs row number (CHIPY) for several cases. At the top, we show Aluminum K-alpha (1.49 keV), while the bottom plot is for Manganese K-alpha (5.9 keV). In each panel, the data points represent the response of the I3 chip (upper curve, without CTI correction; lower curve, with CTI correction), while the lines represent the response of S3. Data were taken from May through July of 2009, on I3 node 3, and S3 node 0 (where the aimpoints are). Nominal aim points are given in chip coordinates in tables 6.4 and 6.5.

Link to postscript file for Figure 6.9 (top)
Link to postscript file for Figure 6.9 (bottom)

The energy resolution (here taken to mean the full width at half maximum [FWHM] of a narrow spectral line) of ACIS varies roughly as the square root of the energy, and increases with distance from the readout. On FI chips (the I array and all the S array chips but S1 and S3), the increase with CHIPY (row number) is dramatic, as can be seen in figure 6.9. The spatial dependence on BI chips (S1 and S3) is much weaker, but depends on both coordinates. To illustrate, we have measured the resolution at the default aimpoints (see table 6.3), using the aluminum (1.49 keV) and manganese (5.9 keV) K-alpha lines in the external calibration source. Data were summed over three months in the summer of 2009, and over an area of 32×32 pixels.

Table 6.3: Spectral Resolution at default aimpoints

ccd	chipx	chipy	1.49 keV	5.9 keV
I3, no offset	937	987	155.757	284.542
I3, w offset	967	963	131.488	256.309
S3, no offset	225	485	95.3451	147.033
S3, w offset	225	516	96.7391	151.302

6.8 Hot Pixels and Columns

Hot pixels and columns are defined to be those which produce a high spurious or saturated pulse-height for a large number of consecutive frames of data. These depend on operating conditions such as temperature. One should always refer to the CXC web site for the most recent list. To date, S1 is the device with the largest number of such pixels and columns.

6.9 Cosmic Ray Afterglows

Cosmic ray hits sometimes deposit so much charge that it appears in a number of successive frames in the same pixels. These cosmic ray afterglows are essentially temporary hot pixels, and are removed by the hot pixel logic if they contain more than approximately 8 events.

6.10 Aimpoints

Aimpoints are the nominal positions on the ACIS where the flux from a point source with no target offsets is placed. There are two nominal aimpoints, indicated in Figure 6.1 - one on the corner of I3 on the ACIS-I array (the ACIS-I aimpoint), and one near the boundary between nodes 0 and 1 on S3 of the ACIS-S array (the ACIS-S aimpoint). Their exact positions are given in Table 6.4. Note that the aimpoint is not the same as the optical axis, which is defined as the position of the narrowest PSF and found approximately 10" from the aimpoints. Because of thermal relaxation of spacecraft structures, the chip coordinates of the optical axis have changed over the mission lifetime. For zero SIM-Z offsets, the aimpoint on ACIS-I is at the I3 chip coordinates (945,983) and on ACIS-S, at the S3 chip coordinates (227,491). Approximate contours of constant encircled energy for ACIS-I and ACIS-S observations for the default aimpoints are shown in Figures 6.10 and 6.11. If required, other aimpoints may be selected along the Z-axis.

Figure 6.10: Approximate contours of constant 50% encircled energy at 1.49 keV when the ACIS-I default aimpoint is selected. The dotted line is 1 arcsec, the dashed line is 1.5 arcsec. The remainder are at 1 arcsec intervals. The thicker solid lines highlight the 5, 10, and 15 arcsec contours.

Link to postscript file for Figure 6.10

Figure 6.11: Approximate contours of constant 50% encircled energy at 1.49 keV when the ACIS-S default aimpoint is selected. The dotted line is 1 arcsec, the dashed line is 1.5 arcsec. The remainder are at 1 arcsec intervals. The thicker solid lines highlight the 5, 10, 15 and 20 arcsec contours.

Link to postscript file for Figure 6.11

In late 2006, a 10^′′ shift in the alignment of the aspect camera occurred when its primary focal plane temperature was cooled from −15C to −19C. In 2011 July, Chandra entered a safemode, and since that time the aimpoint has drifted back to near the 2007 position. However, its short term instability has increased in the past two years or so, especially after the safemode. In combination with the long-term alignment drift, the ACIS aimpoints using the previous default Y and Z offsets have shifted. In order to place the aimpoint so that targets are at the optimal positions with respect to CHIP geometry and telescope focus, we have updated the default offsets for ACIS-S to ∆Y = +9^′′ and ∆Z = −15^′′, to avoid dithering across the node boundary on the S3 chip. For ACIS-I the default offsets of ∆Y = −12^′′ and ∆Z = −15^′′ are unchanged. For grating observations an offset toward the readout on the ACIS-S array is often recommended (i.e. in the negative Z direction - see Table 6.5). Note that standard ACIS subarrays (Section 6.12.1) will not center zeroth order if the recommended aimpoints are selected for a grating observation. In this case, a custom subarray is necessary (e.g. see Section 8.5)

Further information can be found in the memos linked from this page: http://cxc.harvard.edu/cal/Hrma/OpticalAxisAndAimpoint.html

Lastly, it should be kept in mind that the observatory is typically dithered about the aimpoint with an 8^′′ half-amplitude (see Section 6.11, and 6.5 for recommended offsets).

Table 6.4: Aimpoint positions summarized in pixels (chip-x, chip-y)

ACIS-I :	(945, 983) in I3	no offsets
ACIS-I :	(976, 959) in I3	with offsets of ∆Y=−12^′′, ∆Z=−15^′′
ACIS-S :	(227, 491) in S3	no offsets
ACIS-S :	(209, 521) in S3	with default offsets ∆Y=+9^′′, ∆Z=−15^′′

Table 6.5: Recommended SIM-Z offsets

Observation Mode	SIM-Z Offset	Source Position (w/Offsets
Observation Mode	SIM-Z Offset	∆Y=+9^′′, ∆Z=−15^′′)
ACIS-S w/ HETG TE mode:	-3mm = -1.02417′	(209, 396)
ACIS-S w/ HETG CC mode:	-4mm = -1.36556′	(209, 354)
ACIS-S w/ LETG TE mode:	-8mm = -2.73112′	(209, 187)
ACIS-S w/ LETG CC mode:	-8mm = -2.73112′	(209, 187)

Figure 6.12: Position of the optical axis on the S3 CCD at various times in the mission. Since 1999, the aimpoint (position of optimal focus) has drifted (dark lines), and default pointing offsets (lighter lines) have been changed several times to move the image away from the node boundary. Dates are given in format YYYY-MM (year, followed by month).

Link to postscript file for Figure 6.12

6.11 Dither

Unless specially requested, the spacecraft is dithered during all observations. The dither pattern is a Lissajous figure. For observations with ACIS, the dither pattern spans 16 arcsec peak to peak. The dither serves two purposes: (1) to provide some exposure in the gaps between the CCDs, and (2) to smooth out pixel-to-pixel variations in the response. The dither is removed during high-level ground processing of the data. The exposure time in the gaps between chips (and at the outside edges) will be less than that for the remainder of the field. Default dither parameters are listed in Table 5.4.

6.11.1 Gaps Between the CCDs

The approximate sizes of the various gaps between chips are shown in Figure 6.1. Note that the Y-gaps in the ACIS-I array vary with Z due to the way the CCDs are tilted.

6.12 Operating Modes

Here we describe the ACIS operating modes. There are only two such modes. One is Timed Exposure (TE) and the other is Continuous Clocking (CC). One must select one or the other for an observation as it is not possible to simultaneously operate individual CCDs in different modes during a single observation.

6.12.1 Timed Exposure (TE) Mode

A timed exposure refers to the mode of operation wherein a CCD collects data (integrates) for a preselected amount of time - the Frame Time. Once this time interval has passed, the charge from the 1024 x 1024 active region is quickly ( ∼ 41 ms) transferred to the framestore region and subsequently read out through (nominally) 1024 serial registers.

Frame Times - Full Frames

Frame times are selectable within a range of values spanning the time interval from 0.2 to 10.0 seconds. If the data from the entire CCD are utilized (full frame) then the nominal (and optimal!) frame time is 3.2 s, T_opt. Selecting a frame time shorter than the nominal value (e.g. to decrease the probability of pileup - Section 6.15) has the consequence that there will be a time during which no data are taken, "deadtime", as 3.2 s are required for the frame store readout process regardless of the frame time. The fraction of time during which data are taken is simply the ratio of the selected frame time to the sum of the nominal frame time and the 41 msec transfer time - e.g. for a new frame time of t ( < 3.2) secs, the fraction of time during which data are taken is t/(T_opt+0.041). We note, strictly speaking, the full-frame time depends on how many CCD s are on - see the equation below - but the differences are very small. Finally, we note that selecting a frame time longer than the most efficient value increases the probability of pileup occurring and is not recommended.

Frame Times & Subarrays

It is also possible for one to select a subarray - a restricted region of the CCD in which data will be taken. A subarray is fully determined by specifying the number of rows separating the subarray from the framestore region (q) and the number of rows in the subarray (n). Examples of subarrays are shown in Figure 6.13. The nominal frame time for a subarray depends on (q), (n), and the total number of CCDs that are activated (m) - see Table 6.6. The nominal frame time is given by:

T(msec) = 41.12×m +0.040 ×(m×q) + 2.85×n − 32.99.

As with full frames, selecting a frame time less than the most efficient value results in loss of observing efficiency. Frame times are rounded up to the nearest 0.1 sec, and can range from 0.2 to 10.0 sec

When operating with only one chip, subarrays as small as 100 rows are allowed (this permits 0.3 sec frame times which pay no penalty in dead time). For multichip observations, the smallest allowed number of rows is 128.

Table 6.6: CCD Frame Time (seconds) for Standard Subarrays

Subarray	ACIS-I (no. of chips)		ACIS-S (no. of chips)
	1	6	1	6
1	3.0	3.2	3.0	3.2
1/2	1.5	1.8	1.5	1.8
1/4	0.8	1.2	0.8	1.1
1/8	0.5	0.8	0.4	0.7

Figure 6.13: Examples of various subarrays. The heavy dot in the lower left indicates the origin

Link to postscript file for Figure 6.13

Trailed Images

It takes 40 μsec to transfer the charge from one row to another during the process of moving the charge from the active region to the framestore region. This has the interesting consequence that each CCD pixel is exposed, not only to the region of the sky at which the observatory was pointing during the long (frame time) integration, but also, for 40 μsec each, to every other region in the sky along the column in which the pixel in question resides. Figure 6.14 is an example where there are bright features present, so intense, that the tiny contribution of the flux due to trailing is stronger than the direct exposure - hence the trailed image is clearly visible. Trailed images are also referred to as "read out artifact" and "out-of-time images". The user needs to be aware of this phenomenon as it has implications for the data analysis - e.g. estimates of the background. In some cases, the trailed image can be used to obtain an unpiled spectrum and can also be used to perform 40 microsecond timing analysis (of extremely bright sources).

Figure 6.14: Trailed image of a strong X-ray source. The core of the image is faint due to pileup. Most events here are rejected because of bad grades. The readout direction is parallel to the trail.

Link to postscript file for Figure 6.13

6.12.2 Alternating Exposures

In some instances, it is desirable to have both long and short frame times. If the exposure time is made very short, pile-up may be reduced, but the efficiency of the observation is greatly reduced by the need to wait for the full 3.2 seconds (if six chips are clocked) for the framestore array processing.

With alternating exposure times, all CCDs are clocked in unison, but have two exposure times. One (typically short) primary exposure of length 0.2 < t_p < 10.0 sec is followed by k standard exposures t_s (for example, 3.2 seconds if six chips are clocked). Permissible values of k range from zero (the standard single exposure mode) to 15. The short exposures are used to reduce photon pileup, and the long exposures are useful for studying the fainter objects in the field of view. For example, a typical choice of long and short exposure times might be 3.2 and 0.3 seconds. If k = 3, ACIS would perform one 0.3 second exposure, followed by three exposures of 3.2 seconds, repeating until the total observing time expires.

If the duty cycle of long exposures is 1:k (short:long), the observing efficiency η is then

η = 1 −

t_s−t_p

t_s(k+1)

for t_p < 3.2 sec and t_s ≥ 3.2 sec.

6.12.3 Continuous Clocking (CC) Mode

The continuous clocking mode is provided to allow 3 msec timing at the expense of one dimension of spatial resolution. In this mode, one obtains 1 pixel x 1024 pixel images, each with an integration time of 2.85 msec. Details as to the spatial distribution in the columns are lost - other than that the event originated in the sky along the line determined by the length of the column.

In the continuous clocking mode, data is continuously clocked through the CCD and framestore. The instrument software accumulates data into a buffer until a virtual detector of size 1024 columns by 512 rows is filled. The event finding algorithm is applied to the data in this virtual detector and 3 x 3 event islands are located and recorded to telemetry in the usual manner (Section 6.14.1). This procedure has the advantage that the event islands are functionally equivalent to data accumulated in TE mode, hence differences in the calibration are minimal. The row coordinate (called CHIPY in the FITS file) maps into time in that a new row is read from the framestore to the buffer every 2.85 msec. This does have some minor impacts on the data. For example, since the event-finding algorithm is looking for a local maximum, it cannot find events on the edges of the virtual detector. Hence CHIPX cannot be 1 or 1024 (as in TE mode). Moreover, CHIPY cannot be 1 or 512. In other words, events cannot occur at certain times separated by 512*2.85 msec or 1.4592 sec. Likewise, it is impossible for two events to occur in the same column in adjacent time bins.

The TIMEs in continuous-clocking mode Level 1 event data files are the read-out times, not the times of arrival. The differences between the read-out times and the times of arrival are in the range 2.9 to 5.8 s. The differences depend on the nominal location of the source on the CCD and the dither of the telescope. Code to compute the times of arrival at the spacecraft from the read-out times has been developed and was released as part of the tool acis_process_events in CIAO 3.0. Data processed with this version of acis_process_events will have Level 2 files with corrected times.

6.13 Bias Maps

In general, the CCD bias, the amplitude of the charge in each pixel in the absence of external radiation, is determined at various times - every change of instrumental parameters or setup when ACIS is in place at the focus of the telescope. These bias maps have proven to be remarkably stable and are automatically applied in routine data processing.

The bias maps for continuous-clocking mode observations can be corrupted by cosmic rays. If a cosmic ray deposits a lot of charge in most of the pixels in one or more adjacent columns, the bias values assigned to these columns will be too large. As a result, some low-energy events that would have been telemetered will not be telemetered because they do not satisfy the minimum pulse height criterion and the spectrum of a source in the affected columns will be skewed to lower energies. The BI CCDs are relatively insensitive to the problem. A bias algorithm was implemented in Cycle 6 to mitigate the problem.

Occasionally a cosmic ray produces an artifact in a bias map. The pipelines search for these artifacts, and, when found, replace the bias map with another from the same epoch. Work is in progress to use long-term average bias maps, either when there are artifacts in the observation-specific bias map, or for all observations.

6.14 Event Grades and Telemetry Formats

6.14.1 Event Grades

During the first step in the algorithm for detecting X-ray events, the on-board processing examines every pixel in the full CCD image (even in the continuous clocking mode (Section 6.12.3)) and selects as events regions with bias-subtracted pixel values that both exceed the event threshold and are greater than all of the touching or neighboring pixels (i.e., a local maximum). The surrounding 3×3 neighboring pixels are then compared to the bias-subtracted split-event threshold; those that are above the threshold establish the pixel pattern. On the basis of this pattern, the event is assigned a grade. Depending on the grade, the data are then included in the telemetry. On-board suppression of certain grades is used to limit the telemetry bandwidth devoted to background events (see Section 6.16.1).

The grade of an event is thus a code that identifies which pixels, within the 3×3 pixel island centered on the local charge maximum, are above certain amplitude thresholds. The thresholds are listed in Table 6.1. Note that the local maximum threshold differs for the FI and the BI CCDs. A Rosetta Stone to help one understand the ACIS grade assignments is shown in Figure 6.15, and the relationship to the ASCA grading scheme is given in Table 6.7.

Figure 6.15: Schematic for determining the grade of an event. The grade is determined by summing the numbers for those pixels that are above their thresholds. For example, an event that caused all pixels to exceed their threshold is grade 255. A single pixel event is grade 0.

Link to postscript file for Figure 6.15

Table 6.7: ACIS and ASCA Grades

ACIS Grades	ASCA Grade	Description
0	0	Single pixel events
64 65 68 69	2	Vertical Split Up
2 34 130 162	2	Vertical Split Down
16 17 48 49	4	Horizontal Split Right
8 12 136 140	3	Horizontal Split Left
72 76 104 108	6	"L" & Quad, upper left
10 11 138 139	6	"L" & Quad, down left
18 22 50 54	6	"L" & Quad, down right
80 81 208 209	6	"L" & Quad, up right
1 4 5 32 128	1	Diagonal Split
33 36 37 129	1
132 133 160 161	1
164 165	1
3 6 9 20 40	5	"L"-shaped split with corners
96 144 192 13 21	5
35 38 44 52 53	5
97 100 101 131	5
134 137 141 145	5
163 166 168 172	5
176 177 193 196	5
197	5
24	7	3-pixel horizontal split
66	7	3-pixel vertical split
255	7	All pixels
All other grades	7

It is important to understand that most, if not all, calibrations of ACIS are based on a specific subset of ACIS grades. This "standard" set comprises ASCA grades 0,2,3,4, and 6 - G(02346). In the absence of pileup, this particular grade selection appears to optimize the signal-to-background ratio, but this conclusion depends on the detailed spectral properties of the source. Further, most of the scientifically important characteristics of ACIS (effective area, sensitivity, point spread function, energy resolution, etc.) are grade- and energy-dependent.

6.14.2 Telemetry Formats

There are a number of telemetry formats available. Specifying a format determines the type of information that is included in the telemetry stream. The number of bits per event depends on which mode and which format is selected. The number of bits per event, in turn, determines the event rate at which the telemetry will saturate and data will be lost until the on-board buffer empties. The formats available depend on which mode (Timed Exposure or Continuous Clocking) is used. The modes, associated formats, and approximate event rates at which the telemetry saturates and one begins to limit the return of data, are listed in Table 6.8. The formats are described in the following paragraphs. Event "arrival time" is given relative to the beginning of the exposure.

Table 6.8: Telemetry Saturation Limits

Mode	Format	Bits/event	Events/sec^*	Number of Events
				in full buffer
CC	Graded	58	375.0	128,000
CC	Faint	128	170.2	58,099
TE	Graded	58	375.0	128,000
TE	Faint	128	170.2	58,099
TE	Very Faint	320	68.8	23,273

*(includes a 10% overhead for housekeeping data)

Faint

Faint format provides the event position in detector coordinates, an arrival time, an event amplitude, and the amplitude of the signal in each pixel in the 3 x 3 island that determines the event grade. The bias map is telemetered separately. Note that certain grades may be not be included in the data stream (Section 6.16.1).

Graded

Graded format provides event position in detector coordinates, an event amplitude, the arrival time, and the event grade. Note that certain grades may be not be included in the data stream (Section 6.16.1).

Very Faint

Very Faint format provides the event position in detector coordinates, the event amplitude, an arrival time, and the pixel values in a 5 x 5 island. As noted in Table 6.8, this format is only available with the Timed Exposure mode. Events are still graded by the contents of the central 3 x 3 island. Note that certain grades may be not be included in the data stream (Section 6.16.1). This format offers the advantage of reduced background after ground processing (see Section 6.16.2) but only for sources with low counting rates that avoid both telemetry saturation and pulse pileup.

Studies (see http://cxc.harvard.edu/cal/Acis/Cal_prods/vfbkgrnd) of the ACIS background have shown that for weak or extended sources, a significant reduction of background at low and high energies may be made by using the information from 5×5 pixel islands, i.e. very faint mode, instead of the faint mode 3×3 island. This screening results in a 1-2% loss of good events. CIAO 2.2 and later provides a tool to utilize the VF mode for screening background events. Please note that the rmf generation is the same for very faint mode as it is for faint mode. See the "Why Topic" http://cxc.harvard.edu/ciao/why/aciscleanvf.html.

It is important to realize that VF mode uses more telemetry; the limit is ∼ 68 cts/s, which includes the target flux and the full background from all chips. Check the calibration web page (http://cxc.harvard.edu/cal/Acis/Cal_prods/bkgrnd/current/background.html) for a discussion of background flares and the telemetry limit. In particular, review section 1.3 of the memo "General discussion of the quiescent and flare components of the ACIS background" by M. Markevitch

In order to reduce the total background rate and the likelihood of telemetery saturation, VF observations should consider using no more than 5 CCDs and an energy filter with a 12 keV upper cutoff. Starting from Cycle 11, the default upper energy cutoff has been decreased from 15 keV to 13 keV. But VF mode observers should consider reducing this threshold further to 12 keV. If the target is brighter than 5-10 cts/s, one has to take more drastic steps, such as turning off more chips or using Faint mode.

This mode should not be used for observing bright sources (see the discussion at the end of Section 6.16.1 for more detail).

6.15 Pileup

Pileup results when two or more photons are detected as a single event. The fundamental impacts of pileup are: (1) a distortion of the energy spectrum - the apparent energy is approximately the sum of two (or more) energies; and (2) an underestimate as to the correct counting rate - two or more events are counted as one. A simple illustration of the effects of pileup is given in Figure 6.16. There are other, somewhat more subtle, impacts discussed below (6.15.1).

Figure 6.16: The effects of pileup at 1.49 keV (Al Kα) as a function of source intensity. Data were taken during HRMA-ACIS system level calibration at the XRCF. Single-photon events are concentrated near the pulse height corresponding to the Al Kα line ( ∼ 380 ADU), and events with 2 or more photons appear at integral multiples of the line energy.

Link to postscript file for Figure 6.16

The degree to which a source will be piled can be roughly estimated using PIMMS. Somewhat more quantitative estimates can be obtained using the pileup models in XSPEC, Sherpa and ISIS. If the resulting degree of pileup appears to be unacceptable given the objectives, then the proposer should employ some form of pileup mitigation (Section 6.15.3) as part of the observing strategy. In general, pileup should not be a problem in the observation of extended objects, the Crab Nebula being a notable exception, unless the source has bright knots or filaments.

6.15.1 Other Consequences of Pileup

There are other consequences of pileup in addition to the two principal features of spurious spectral hardening and underestimating the true counting rate by under counting multiple events. These additional effects are grade migration and pulse saturation, both of which can cause distortion of the apparent PSF.

Grade migration

Possibly the most troubling effect of pileup is that the nominal grade distribution that one expects for X-ray events changes. The change of grade introduced by pileup has become to be referred to as "grade migration". Table 6.9 shows an example of grade migration due to pileup as the incident flux is increased. In this simple test, which involved only mono-energetic photons, the largest effect is the depletion of G0 events and the increase of G7 events. In general, as the incident flux rate increases, the fraction of the total number of events occupying a particular event grade changes as photon-induced charge clouds merge and the resulting detected events "migrate" to other grades which are not at all necessarily included in the standard (G02346) set. If one then applies the standard calibration to such data, the true flux will be under-estimated.

Table 6.9: ASCA Grade Distributions for different incident fluxes at 1.49 keV (Al-Kα, based on data taken at the XRCF during ground calibration using chip I3)

Incident
Flux^*	G0	G1	G2	G3	G4	G5	G6	G7
9	0.710	0.022	0.122	0.053	0.026	0.009	0.024	0.035
30	0.581	0.057	0.132	0.045	0.043	0.039	0.029	0.073
98	0.416	0.097	0.127	0.052	0.050	0.085	0.064	0.108
184	0.333	0.091	0.105	0.040	0.032	0.099	0.077	0.224

^*arbitrary units

Pulse Saturation

One consequence of severe instances of pileup is the creation of a region with no events! In this case, the pileup is severe enough that the total amplitude of the event is larger than the on-board threshold (typically 13 keV) and is rejected. Holes in the image can also be created by grade migration of events into ACIS grades (e.g. 255) that are filtered on-board.

Figure 6.17: The effects of pileup on the radial distribution of the PSF are illustrated. These data were taken during ground calibration at the XRCF. The specific "OBSIDs", the counting rate per CCD frame ("c/f"), and the "pileup fraction" as defined in Section 6.15.2 are given in the inset.

Link to postscript file for Figure 6.17

PSF distortion

Obviously the effects of pileup are most severe when the flux is highly concentrated on the detector. Thus, the core of the PSF suffers more from pileup induced effects than the wings. Figure 6.17 illustrates this point.

6.15.2 Pileup Estimation

It is clearly important in preparing a Chandra observing proposal to determine if the observation will be impacted by pileup, and if so, decide what to do about it (or convince the peer review why the specific objective can be accomplished without doing anything). There are two approaches to estimating the impact of pileup on the investigation. The most sophisticated uses the pileup models in XSPEC, Sherpa, and ISIS to create a simulated data set which can be analyzed in the same way as real data. A less sophisticated, but very useful, approach is to use the web version of PIMMS to estimate pileup or to use Figures 6.17 and 6.18.

Figure 6.18: The pileup fraction as a function of the the counting rate (in the absence of pileup in units of photons/frame). The solid line is for on-orbit, the dashed line and the data points are for, and from, ground-based data respectively. The difference between the ground and flight functions are a consequence of the improved PSF on-orbit, where gravitational effects are negligible - see Chapter 4. Note that when pileup occurs there are two or more photons for each event, so the fraction of events with pileup is always less than the fraction of photons with pileup.

Link to postscript file for Figure 6.18

Simple Pileup Estimates The pileup fraction is the ratio of the number of detected events that consist of more than one photon to the total number of detected events. An estimate of the pileup fraction can be determined from Figure 6.18. The algorithm parameterizes the HRMA+ACIS PSF in terms of the fraction of encircled energy that falls within the central 3×3 pixel event detection cell, and assumes that the remaining energy is uniformly distributed among the 8 surrounding 3×3 pixel detection cells. The probabilities of single- and multiple-photon events are calculated separately for the central and surrounding detection cells and subsequently averaged (with appropriate weighting) to obtain the pileup fraction as a function of the true count rate - the solid line in Figure 6.18. The model was tested against data taken on the ground under controlled conditions - also shown in Figure 6.18.

Figure 6.19: MARX simulations of the effect of pileup on the shape of the spectrum. The true (solid line) and the detected (dotted line) spectra are shown for four different viewing angles. The corresponding "pileup fractions" - see Section 6.15.2 - are 46%, 40%, 15%, and 2% as the image is moved progressively further off-axis.

LINK TO POSTSCRIPT FILE FOR Figure 6.19

As a general guideline, if the estimated pileup fraction is > 10%, the proposed observation is very likely to be impacted. The first panel (upper left) in Figure 6.19 qualitatively illustrated the effect on a simulated astrophysical X-ray spectrum. However, the degree of pileup that is acceptable for a particular objective will depend on the particular scientific goals of the measurement, and there is no clear-cut tolerance level. If one's scientific objective demands precise flux calibration, then the pileup fraction should probably be kept well below 10%.

The PIMMS tool provides the pileup fraction using the algorithm described here, both for direct observation with ACIS and also for the zeroth-order image when a grating is inserted.

Simulating Pileup

John Davis at MIT has developed an algorithm for modeling the effects of pileup on ACIS spectral data. The algorithm has been implemented as of XSPEC V11.1 and Sherpa V2.2. The algorithm can be used to attempt to recover the underlying spectrum from a source, or to simulate the effects of pileup for proposal purposes.

The algorithm has been tested by comparing CCD spectra with grating spectra of the same sources. Care should be taken in applying the algorithm - for example, using the appropriate regions for extracting source photons and avoiding line-dominated sources. A description of the algorithm can be found in Davis 2001 (Davis, J.E. 2001, ApJ, 562, 575). Details on using the algorithm in Sherpa are given in a Sherpa "thread" as of CIAO V2.2 on the CXC CIAO web page: http://cxc.harvard.edu/ciao/.

6.15.3 Reducing Pileup

We summarize here various methods which can be used to reduce pileup.

•Shorten exposure time:: By cutting back on CCD exposure time, the probability of pileup decreases. The user is advised to select the best combination of a subarray and frame time in order to avoid losing data as discussed in Section 6.12.1.
•Use the Alternating Exposure option:: This option simply alternates between exposures that are subject to pileup and those that are not. The capability was originally developed for use with certain grating observations to allow one to spend some time obtaining useful data from a zeroth order image, which would otherwise be piled up. See section 6.12.2.
•Use CC mode: If two-dimensional imaging is not required, consider using CC mode (Section 6.12.3).
•Insert a transmission grating:: Inserting either the HETG (Chapter 8) or the LETG (Chapter 9) will significantly decrease the counting rate as the efficiency is lower. The counting rate in the zero order image may then be low enough to avoid pileup.
•Offset point:: Performing the observation with the source off-axis spreads out the flux and thus decreases the probability of pileup at the price of a degraded image. Figure 6.19 illustrated the impact.
•Defocus:: The option is only listed for completeness, the option is not recommended or encouraged.

6.16 On-Orbit Background

There are three components to the on-orbit background. The first is that due to the cosmic X-ray background (a significant fraction of which resolves into discrete sources during an observation with Chandra). The second component is commonly referred to as the charged particle background. This latter arises both from charged particles, photons, and other neutral particle interactions that ultimately deposit energy in the instrument. The third component is the "readout artifact" which is a consequence of the "trailing" of the target image during the CCD readout; it is discussed in Section 6.12.1.

The background rates differ between the BI and the FI chips, in part because of differences in the efficiency for identifying charged particle interactions. Figure 6.20 illustrates why.

Figure 6.20: Enlarged view of an area of a FI chip I3 (left) and a BI chip (right) after being struck by a charged particle. There is far more "blooming" in the FI image since the chip is thicker. The overlaid 3×3 detection cells indicate that the particle impact on the FI chip produced a number of events, most of which end up as ASCA Grade 7, and are thus rejected with high efficiency. The equivalent event in the BI chip, is much more difficult to distinguish from an ordinary X-ray interaction, and hence the rejection efficiency is lower.

Link to postscript file for Figure 6.20 (left) Link to postscript file for Figure 6.20 (right)

6.16.1 The Non-X-ray Background

Starting from September 2002, several observations were carried out with the ACIS in the stowed position, shielded from the sky by the SIM structure, and collecting data in normal imaging TE VF mode at -120C. Chips I0, I2, I3, S1, S2, S3 were exposed. The SIM position was chosen so that the on-board calibration source did not illuminate the ACIS chips. This allowed us to characterize the non-celestial contribution to X-ray background (i.e., from charged particles). The resulting spectra from different chips are shown in Figure 6.21. Chip S2 is similar to the ACIS-I chips (denoted I023 in the figure) and not shown for clarity.

In addition, in July-September 2001, Chandra performed several short observations of the dark Moon, which blocks the cosmic X-ray background. The dark Moon and stowed background spectra were indistinguishable (except for short periods of flares and variable Oxygen line emission in the Moon observations). We have not observed any background flares in the stowed position. Thus, the ACIS-stowed background is a good representation of the quiescent non-X-ray background in the normal focal position and can be used for science observations.

Figure 6.21: Energy spectra of the charged particle ACIS background with ACIS in the stowed position (a 50 ks exposure taken in September 2002; standard grade filtering, no VF filtering). Line features are due to fluorescence of material in the telescope and focal plane.

Link to postscript file for Figure 6.21

The flight grade distributions in early measurements of the non-X-ray background for the two types of CCDs are shown in Figure 6.22. Although subsequent to these early measurements the CCD temperature has been lowered and the FI devices suffered the effects of the radiation damage, the background is still dominated by the same grades. Based on these data, events from flight grades 24, 66, 107, 214, and 255 are routinely discarded on-board. The total rate of the discarded events is available in the data stream. The remaining non-X-ray events telemetered to the ground are still dominated (70-95%) by other bad grades. They are not discarded onboard because some of them may turn out to be valid X-ray events after ground processing.

For data taken using the Very Faint (VF) telemetry format (Section 6.14.2), the non-X-ray background can be reduced in data processing by screening out events with significant flux in border pixels of the 5×5 event islands. This screening leaves the data from faint sources essentially the same while reducing the FI background at different energies: a factor of 1.4 (E > 6 keV); 1.1 (1-5 keV); and 2 (near 0.5 keV). For the BI chips the reductions are: 1.25 (E > 6 keV); 1.1 (1-5 keV); and 3 (near 0.3 keV). The screening algorithm has been incorporated into the CIAO tool acis_process_events. Further discussion may be found at http://cxc.harvard.edu/cal/Acis/Cal_prods/vfbkgrnd/index.html

Proposers should be aware that telemetry saturation occurs at lower count rates for observations using the VF format, so they may need to take steps to limit the total ACIS count rate (see Sec. 6.16.2). Proposers should also be aware that if there are bright point sources in the field of view, the screening criterion discussed above is more likely to remove source events due to pileup of the 5×5 pixel event islands. Point sources should have count rates significantly less than 0.1 counts/sec to be unaffected. However, there is no intrinsic increase of pileup in VF data compared to Faint mode, and the screening software can be selectively applied to regions, excluding bright point-like sources. Thus the use of VF mode is encouraged whenever possible.

Figure 6.22: Fraction of ACIS background events as a function of grade from early in-flight data for an FI chip (S2) (top) and a BI chip (S3) (bottom).

Link to postscript file for Figure 6.22 (top)
Link to postscript file for Figure 6.22 (bottom)

6.16.2 The Total Background

In real observations, two more components to the background come into play. The first is the cosmic X-ray background which, for moderately long ( ∼ 100 ks) observations will be mostly resolved into discrete sources (except for the diffuse component below 1 keV) but, nevertheless, contributes to the overall counting rate. The second is a time-variable "flare" component caused by any charged particles that may reflect from the telescope and have sufficient momentum so as not to be diverted from the focal plane by the magnets included in the observatory for that purpose, or from secondary particles (Section 6.16.3). Figure 6.23 compares flare-free ACIS-S3 spectra of the non-X-ray (dark Moon) background and a relatively deep pointing to a typical region of the sky away from bright Galactic features.

Figure 6.23: ACIS -S3 spectrum of the non-X-ray background (large crosses) overlaid on the quiescent blank sky spectrum. Small crosses show the total sky spectrum, while squares show the diffuse component left after the exclusion of all point sources detectable in this 90 ks exposure. Standard grade filtering and VF filtering are applied. The background and blank-sky spectra are normalized to the same flux in the 10-12 keV band.

Link to postscript file for Figure 6.23

The 2007-2008 quiescent detector+sky background counting rates in various energy bands and for the standard good grades are given in Table 6.10. Insertion of the gratings makes little measurable difference in the background rates, but it does block the background flares. The lower-energy rates are very approximate as they vary across the sky. The rates are slowly changing on the timescale of months, so Table 6.10 can only be used for rough sensitivity estimates. Table 6.11 gives total background count rates for each type of chip, including all grades that are telemetered (see Section 6.14.1 and 6.16.1), and can aid in estimating the probability of telemetry saturation.

	Bkgrd rates (cts/s/chip)
Energy
Band (keV)	I0	I1	I2	I3	S1	S2	S3	S4
0.3-10	0.43	0.49	0.43	0.49	2.26	0.46	1.18	0.54
0.5-2	0.10	0.10	0.11	0.10	0.30	0.07	0.11	0.16
0.5-7	0.26	0.26	0.26	0.27	0.70	0.27	0.51	0.33
5.0-10	0.22	0.22	0.21	0.22	0.96	1.54	0.67	0.24
10-12	0.13	0.13	0.13	0.13	1.15	0.13	0.92	0.14

Table 6.10: Approximate on-orbit standard grade background counting rates. The rates are cts/s/chip, using only ASCA grades 02346, no VF filtering, excluding background flares and bad pixels/columns and celestial sources identifiable by eye. These values can be used for sensitivity calculations

Period	Aug 1999	2000-2003	2009	2011
Upper E cutoff	15 keV	15 keV	15 keV	15 keV	13 keV	12 keV	10 keV
Chip S2 (FI)	10	6.3	10.7	8.1	6.7	6.5	6.1
Chip S3 (BI)	11	7.7	14.7	11.7	8.0	7.2	5.1

Table 6.11: Typical total quiescent background rates (cts/s/chip), including all grades that are telemetered (not just standard ASCA grades), by chip type and upper energy cutoff. These values can be used to estimate the probability of telemetry saturation.

The background rates have declined in 1999-2000, stayed flat through end of 2003, increased until the end of 2009, and have recently started to drop, following the solar cycle. (Fig. 6.24). By late 2011, the rates were roughly halfway between the peak of 2009 and the low rates of 2001-2003. If the background rate is anti-correlated with the solar cycle, we may expect the sharp decline of the rates to continue in 2012, but no assumptions of this can be made for planning of observations.

Figure 6.24: Total telemetered background rates (including all grades and the upper event cutoff at 15 kev) for chips S2 (FI) and S3 (BI) as a function of time. Vertical dashed lines are year boundaries.

Link to postscript file for Figure 6.24

For aid in data analysis and planning background-critical observations, the CXC has combined a number of deep, source-free, flare-free exposures (including all components of the background) into background event files for different time periods. These blank-sky datasets, along with the detector-only (ACIS-stowed) background files (section 6.16.1), can be found in the CALDB.

6.16.3 Background Variability

In general, the background counting rates are stable during an observation. Furthermore, the spectral shape of the non-X-ray background has been remarkably constant during 2000-2009 for FI and, to a lower accuracy, for BI chips, even though the overall background rate showed secular changes by a factor of 1.5. (For chip S3, the shape has been constant during 2000-2005, but a small change is being observed since late 2005.) When the quiescent background spectra from different observations are normalized to the same rate in the 10-12 keV interval, they match each other to within ±3% across the whole Chandra energy band. The previous discussion assumes that the upper threshold will be set to 13 keV.

Occasionally, however, there are large variations (flares), as illustrated in Figure 6.25. Figure 6.26 shows the frequency of such variations when compared to the quiescent background. An average fraction of the exposure affected by flares above the filtering threshold used for the blank-sky background datasets (a factor of 1.2 above the nominal rate) was about 6% for FI chips and up to 1/3 for BI chips during the first few years of the mission. This fraction has declined with time, and is practically zero at present (mid-2010). If the background follows the solar cycle, the flare frequency may soon return to its early values. Thus, given that the quiescent background in FI chips is also lower than that in S3, background-critical observations may best be done with ACIS-I.

Several types of flares have been identified, including flares that are seen only in the BI chips, and flares that are seen in both the FI and BI chips. Figure 6.27 shows the spectra of two of the most common flare species. Both flares have spectra significantly different from the quiescent background. In addition, the BI flares have spatial distribution very different from that of the quiescent background. The BI flares produce the same spectra in S1 and S3.

Users should note that the total counting rate can significantly increase during a flare (although flare events are almost exclusively good-grade so the total rate does not increase by as large a factor as the good-grade rate; details can be found at http://cxc.harvard.edu/cal/Acis/Cal_prods/bkgrnd/current). If the probability of telemetry saturation is significant, users of ACIS-I might consider turning off the S3 chip. However, if ACIS-S is used in imaging mode, the CXC recommends that both BI chips be turned on. The advantage is that for most types of flares S1 can be used to determine the flare spectrum which can then be subtracted from the spectrum obtained with S3.

Figure 6.25: An example of the ACIS background counting rate versus time - BI chip (S3; top curve) and an FI chip (I2; bottom curve). These are for the standard grades and the band from 0.3 - 10 keV.

Link to postscript file for Figure 6.25

Figure 6.26: An estimate of the cumulative probability that the ratio of the background counting rate to the quiescent background counting rate is larger than a given value. Upper plot for a representative FI chip - S2, and the lower curve for a representative BI chip - S3. The vertical dotted line is a limiting factor 1.2 used in creating the background data sets. These probabilities are relevant for the archival data from the first 2-3 years of the mission. They declined with time to almost zero at present, but may soon increase again as we enter the new solar cycle.

Link to postscript file for Figure 6.26

Figure 6.27: Spectra of different background flares in chip S3. Thick crosses show a common flare species that affects only the BI chips. Thin crosses show one of the several less common flare species that affect both the BI and FI chips. Note how both these spectra are different from the quiescent spectrum (see Figure 6.21).

Link to postscript file for Figure 6.27

6.16.4 Background in Continuous Clocking Mode

Apart from compressing the data into one dimension (Section 6.12.3), there is essentially no difference in the total background in CC mode and that encountered in the timed exposure mode. The background per-sky-pixel, however, will be 1024 times larger, since the sky-pixel is now 1 x 1024 ACIS pixels.

6.17 Sensitivity

The sensitivity for detecting objects is best estimated using the various proposal tools such as PIMMS, MARX, etc. The "Chandra Proposal Threads" web page gives detailed examples of how to use these tools (http://cxc.harvard.edu/proposer/threads)

6.18 Bright Source X-ray Photon Dose Limitations

Pre-Flight radiation tests have shown that ∼ 200 krads of X-ray photon dose will damage the CCDs. The mechanism for the damage is the trapped ionization in the dielectric silicon oxide and nitride separating the gates from the depletion region. Since the charge is trapped, the damage is cumulative. Because the structure of the BI CCDs differs significantly from that of the FI CCDs, the two types of chips have different photon dose limitations. Specifically, the BI CCDs are more than 25 times as tolerant to a dose of X-ray photons as compared to the FI CCDs since the former have ∼ 40 μm of bulk Si protecting the gate layer.

Simulations of astrophysical sources have yielded a very conservative, spectrally-averaged, correspondence of 100 counts/pix = 1 rad. By `counts' in this context we mean all photons that impinged on the detector, whether or not they were piled-up or discarded onboard.

In consultation with the IPI team the CXC has adopted the following mission allowances, per pixel of the two types of chips:

        FI chips: 25 krads   2,500,000 cts/pix
        BI chips: 625 krads  62,500,000 cts/pix

If your observation calls for observing a bright point-like source close to on-axis, we suggest you use the MARX simulator (with the parameter DetIdeal=yes & dither, typically, on) to calculate whether your observation may reach 1% of the above mission limits in any one pixel. If so, please contact the CXC HelpDesk in order to custom design an observational strategy which may accommodate your science aims, while maintaining the health & safety of the instrument.

6.19 Observing Planetary and Solar System Objects with ACIS

Chandra has successfully observed several solar system objects, including Venus, the Moon, Mars, Jupiter and several comets. Observations of planets and other solar system objects are complicated because these objects move across the celestial sphere during an observation and the optical light from the source can produce a significant amount of charge on the detectors (this is primarily an issue for ACIS-S observations). Some information regarding observation planning and data processing is given here. Users are encouraged to contact the CXC for more detailed help.

6.19.1 Observations with ACIS-I

Any solar system object can be observed with ACIS-I. Previous solar system observations with ACIS-I have not revealed significant contamination from optical light. However, proposers are encouraged to work with the CXC when planning the specifics of a given observation. Since the source moves across the celestial sphere in time, an image of the event data will exhibit a "streak" associated with the source. The CIAO tool sso_freeze can be used to produce an event data file with the motion of the source removed.

6.19.2 Observations with ACIS-S

The ACIS-S array can be used with or without a grating. The BI CCDs are more sensitive to soft X-rays than the I array CCDs, but the entire S array suffers from the disadvantage that its OBF is thinner than for ACIS-I and may transmit a non-negligible flux of visible light onto the CCDs. It is thus necessary to estimate the amount of charge produced in the CCDs due to the optical light. More detailed information can be found at http://www.astro.psu.edu/xray/docs/cal_report/ and from the CXC via HelpDesk.

If the optical light leak is small enough, it can be mitigated by simply shortening the frame time. This leads to a linear drop in the number of ADU due to optical light. If possible, VF mode should be used, since in this mode the outer 16 pixels of the 5×5 region allows a "local" bias to be subtracted from the event to correct for any possible light leakage. However, see the warnings in Section 6.14.2.

The optical light also invalidates the bias taken at the beginning of the observation if a bright planet is in the field. It is therefore desirable to take a bias frame with the source out of the field of view. This bias map is useful even when processing 5×5 pixels in VF mode since it can be employed as a correction to the local average "bias" computed from the 16 outer pixels, thereby correcting for hot pixels, cosmetic defects etc.

A more sophisticated approach to dealing with excess charge due to optical light is to make an adjustment to the event and split thresholds. Event grades are described in more detail in Section 6.14.1. Excess charge (in ADU) due to optical light will be added to the event and split counters on-board. Without an adjustment to the thresholds (or a large enough threshold), many of the X-ray events may have all nine pixels of a 3 ×3 pixel event detection cell above the split threshold, in which case the event will not be telemetered to the ground. If the adjustment is too large, X-ray events may not be detected because they may not exceed the event threshold.

Users should be aware that if the detection thresholds are adjusted, standard CXC processing of planetary data will give inaccurate estimates of event pulse heights and grades. To analyze such data, a thorough understanding of the energy calibration process and manual massaging of the data will be required.

6.20 Observing with ACIS - the Input Parameters

This section describes the various inputs that either must be, or can be, specified in order to perform observations with ACIS. The sub-sections are organized to match the RPS form. We have added some discussion as to some of the implications of the possible choices. As emphasized at the beginning of the Chapter, ACIS is moderately complex and the specific characteristics of the CCDs and their configuration in the instrument lead to a number of alternatives for accomplishing a specific objective - detailed trade-offs are the responsibility of the observer. Thus, e.g. it might seem obvious that observations of a faint point source may be best accomplished by selecting the ACIS-S array with the aim point on S3, the BI device that can be placed at the best focus of the telescope, and the CCD with the best average energy resolution. On the other hand, perhaps the science is better served by offset pointing (by a few arcminutes) the target onto S2, very near to the framestore, where the FI energy resolution is better than that of S3. On the other hand, if the object is very faint, so that the number total number of photons expected is just a handful - not enough to perform any significant spectroscopy - the advantage of S3 may not be so obvious considering the smaller field of view, and perhaps the ACIS-I array, which would optimize the angular resolution over a larger field, may be more attractive.

6.20.1 Required Parameters

There are certain ACIS input parameters that must be specified: the number and identity of the CCDs to be used, the Exposure Mode, and the Event Telemetry Format. If pileup and telemetry saturation are not considered to be a problem for the observation, then these are the only parameters that need to be specified.

Number and Choice of CCDs Up to six CCDs can be operated at once, if the science objectives require 6 CCDs. Previously in the mission, the CXC advised GOs to use 6 CCDs if possible. Given the changes in the thermal environment on the spacecraft around ACIS, the CXC now recommends that GOs use 5 or fewer CCDs if there is no impact on their science. Please see the discussion under "Choosing Optional CCDs" that follows for the reasons why we wish the observer to request fewer than 6 if feasible. The identity of the CCDs and the desired aimpoint must be specified.

Choosing Optional CCDs

The observer may specify that a given CCD must be on for an observation by entering "Y" for that CCD at the appropriate place in the RPS form. If a CCD must be off, one enters an "N". Finally, one may specify rank-ordered optional CCDs which will be turned on or off at the discretion of the mission schedulers depending on the specific thermal status of the Observatory at the time of the observation. This is done by entering "OPT1-OPT5" ("O1-O5" in RPS). The first to be turned off, if necessary, would be designated by "OPT1", the second to be turned off would be designated by "OPT2", etc. Observers are encouraged to use 5 or fewer CCDs if their science objectives are not significantly affected. Starting in Cycle 14, the RPS forms for proposal submission will no longer allow the proposer to specify "Y" for 6 CCDs. If the observer's science objectives nevertheless require 6 CCDs to be on, the observer should set 5 CCDs to "Y" and 1 CCD to "OPT1". If the proposal is selected, the observer should discuss the configuration with the user uplink support scientist and the optional CCD may be turned on. The CXC will do its best to schedule 6 CCD observations, but proposers should be aware that satisfying the thermal requirements of the Observatory may make the scheduling of such observations more complicated if not impossible.

Using fewer than 6 CCDs is beneficial in keeping the ACIS focal plane and electronics temperatures within the required operating ranges. Because of changes in the Chandra thermal environment, the ACIS Power Supply and Mechanism Controller (PSMC) has been steadily warming over the course of the mission. Under current conditions, and assuming an initial PSMC temperature of less than +30 C, observations at pitch angles less than 60 degrees, longer than 50 ks, and with 6 CCDs operating, are likely to cause the PSMC to approach or exceed the thermal limit. For pitch angles larger than 130 degrees, the ACIS focal plane, the Detector Electronics Assembly (DEA), and the Digital Processing Assembly (DPA) temperatures can also warm outside of the desired range. As above, each of these temperatures can be reduced by reducing the number of operating CCDs.

If an observer requires the most accurate gain calibration for their observation (provided by a cold and stable focal plane temperature), they should use 5 or fewer CCDs. This is done by setting 5 CCDs to "N", and 5 CCDs to "Y" or 1 CCD to "Y" and 4 other CCDs to "OPT1-OPT4". For example, if the observer is using the ACIS-I array for imaging, they could select the four I array CCDs and one of S2 or S3. If the observer is using S3 for imaging, they could select S2, S3, S4, I2 and I3 and turn S1 off or they could select just S1, S2, S3, and S4.

If no optional CCDs are selected, and six CCDs are on and the observation is not constrained in such a way as to prohibit it, the observation is likely to be scheduled at a time for which the pitch angle is between 60 degrees and 130 degrees. The user should be aware that such an observation may be listed in the long term schedule (LTS) at a date for which the pitch angle for this target is less than 60 degrees or greater than 130 degrees. However, when the observation finally appears in the short term schedule, it will be at a date for which the pitch angle will be in the allowed range.

Some Recommended Chip Sets
Observers should specify the chip set that is best for their primary science. The following suggestions have proven to be popular, and would facilitate a more useful and homogeneous archive.


  ACIS-I (5 CCDs) imaging, nominal aimpoint
            -------
                 I0    I1
                 Y     Y
                 I2    I3
                 Y     Y
       S0    S1  S2    S3   S4   S5
       N     N   Y     N    N    N




  ACIS-I (5 CCDs) imaging, nominal aimpoint
            -------
                 I0    I1
                 Y     Y
                 I2    I3
                 Y     Y
       S0    S1  S2    S3   S4   S5
       N     N   N     Y    N    N




  ACIS-I (6 CCDs) imaging, nominal aimpoint
            -------
                 I0    I1
                 Y     Y
                 I2    I3
                 Y     Y
       S0    S1  S2    S3   S4   S5
       N     N   O2    O1   N    N

The rationale for the first configuration above is that, in the unlikely event of major background flares, telemetry might saturate more rapidly if S3 were on and the FP temperature and electronics temperatures will be lower with only 5 CCDs on. For the middle configuration, the rationale is that S3 is generally more sensitive and closer to the ACIS-I aimpoint, so more sensitive to serendipitous source detection. The rationale for the third configuration is that it is desired to have both S2 and S3 on, but it is not required. Given the current thermal performance of the spacecraft, it is probable that one of the two optional CCDs would be turned off.


ACIS-S (5 CCDs) imaging, nominal aimpoint
            -------
                 I0    I1
                 N     N
                 I2    I3
                 Y     Y
       S0    S1  S2    S3   S4   S5
       N     N   Y     Y    Y    N




  ACIS-S (5 CCDs) imaging, nominal aimpoint
            -------
                 I0    I1
                 N     N
                 I2    I3
                 N     Y
       S0    S1  S2    S3   S4   S5
       N     Y   Y     Y    Y    N




  ACIS-S (6 CCDs) imaging, nominal aimpoint
            -------
                 I0    I1
                 N     N
                 I2    I3
                 O1    O2
       S0    S1  S2    S3   S4   S5
       N     Y   Y     Y    Y    N

For this ACIS-S imaging chipset, chips farthest from the aimpoint (where the PSF is degraded) would be turned off first. On the other hand, S4 has significant noise streaks and resulting decreased sensitivity, so some users may prefer to turn it off earlier in the optional sequence. The rationale for the selection on the first set above is that S1 will have a higher count rate than a FI CCD in the case of a background flare and thus it might be desirable to have S1 off. The rational for the middle selection above is that the observer may want to use S1 help model the background on S3, therefore it is desirable to have I2 turned off. The rational for the third configuration is that it is desirable to have 6 CCDs on, but not required. Given the current thermal performance of the spacecraft, it is probable that one of the two optional CCDs would be turned off.


  ACIS-S (6 CCDs) spectroscopy, nominal aimpoint
            -------
                 I0    I1
                 N     N
                 I2    I3
                 N     N
       S0    S1  S2    S3   S4   S5
       O1     Y   Y     Y    Y    O2

The optimum ACIS-S spectroscopy chip set depends strongly on the expected spectrum of the target. Typically the maximum signal is desired, so the HETG and LETG observer is most likely to insist on all ACIS-S chips. If the science does depend strongly on the flux received on the S0 and S5 CCDs, the observer may need to specify all 6 ACIS-S CCDs to be on. However, the observer should be aware that the amount of useful flux on S0 and S5 with the LETG is typically quite low given they are both FI CCDs (see figure 9.6). Given the current thermal performance of the spacecraft, it is probable that one of the two optional CCDs would be turned off.

ACIS-I Count Rate and Number of Spectral Lines

We want to know the maximum number of counts from the brightest source in the field (not necessarily the target), and if the proposer believes that there will be lines in the spectrum from this source. Both questions are being asked to trigger assessment by the CXC of the sensitivity of the experiment to the gain calibration and so that appropriate scheduling may be employed to avoid thermally-induced gain drifts that might impact the observation. If the GO is interested in the optimal spectral performance of the ACIS CCDs, they should seriously consider using only 4 CCDs as discussed above.

Exposure Mode

There are only two choices: Timed Exposure (Section 6.12.1) or Continuous Clocking (Section 6.12.3).

Timed Exposure Mode

The timed exposure mode with the default nominal (and optimal) frame time of 3.2s is the typical mode for ACIS observations. Note that the option of selecting frame times shorter than nominal reduces observing efficiency, and hence the number of photons collected for a given observation time.

Continuous Clocking Mode

The Continuous Clocking mode is useful when timing data are so critical and/or pileup is such a problem that the sacrifice of one dimension of spatial data is warranted. The use of continuous clocking may also lead one to consider specifying a particular satellite roll orientation (see Chapter 3) in order to avoid having two different sources produce events in the same CCD column. (See section 6.20.4 below.)

6.20.2 Optional Parameters

Alternating Exposures

This option applies only to Timed Exposures. The parameters specifying an alternating exposure are:

the number of secondary exposures per primary exposure (1-15)
the primary exposure frame time

Frame times and efficiencies in TE mode are discussed in Sections 6.12.1 and 6.12.1. The Alternating Exposure option is discussed in Section 6.12.2.

Energy Filtering

It is possible to remove events from the telemetry stream, and thus avoid telemetry saturation, by specifying an energy acceptance filter within which detected events will be telemetered. The default discards events above 3250 ADU (nominally 13 keV). The total per-chip background rates for different upper energy cut-offs are in Table 6.11.

Starting September 2006, a new energy to PHA conversion is used for observations with energy filters. Two sets of conversions are used, depending on the aimpoint of the observation. Observations with ACIS-S at the aimpoint use a conversion tailored for the BI CCDs and those with ACIS-I at the aimpoint use FI CCD-specific conversions. The BI and FI specific conversions are more accurate for each type of CCD than the conversion used in previous cycles. The assumption is that it is desirable to have the most accurate gain conversion for the CCD on which the HRMA aimpoint falls. Note that the conversion only impacts the on-orbit energy filtering. Ground data processing will always apply the appropriate PHA to energy conversion.

The observer should be aware that for observations which mix CCD type (i.e. BI and FI CCDs on), the selected conversion (based on aimpoint as above), will nevertheless apply to all selected CCDs. This will not affect the observation if the low energy threshold for the energy filter (the "Event filter: Lower" parameter) is 0.5 keV or less as the use of either conversion at these energies results in essentially no difference in the number of accepted events. However, above 0.5 keV, the conversions are significantly different. Finally, observations which apply an energy filter with a low energy threshold greater than 0.5 keV will automatically be assigned spatial windows that allow the FI CCDs to use the FI conversion and the BI CCDs to use the BI conversion regardless of aimpoint. Proposers who need an energy filter above 0.5 keV are encouraged to contact the CXC HelpDesk to discuss their plans with an instrument scientist.

Spatial Windows

A more sophisticated approach to removing data from the telemetry stream, and thus avoiding telemetry saturation, is by the use of a Spatial Window. This option offers a good deal of flexibility. One may define up to 6 Spatial Windows per CCD. Each window can be placed anywhere on the chip. Note there is a significant difference between a Spatial Window and a Subarray (Section 6.12.1): Subarrays affect the transmission of CCD data to the on-board ACIS processors; Spatial Windows select events detected by the processors and only impact the telemetry rate. The user may also specify the window energy threshold and energy range.

Spatial windows can specify the sample rate for events inside them. A sample rate of 0 excludes all events; the default rate of 1 includes all events; a rate of n > 1 telemeters one out of every n events in the window. A spatial window could be used to eliminate a bright, off-axis source that would otherwise overwhelm the telemetry stream. The order in which the spatial windows is specified is important if they overlap. The earliest specified window including a given pixel will be applied to events at that pixel.

6.20.3 Non-ACIS Parameters Relevant to an Observation with ACIS

There are a small number of additional parameters that need to be considered in specifying an observation with ACIS: (1) the off-axis pointing (if required), which reduces the flux, and spreads out the image; (2) the roll angle (Chapter 4); (3) time constraints (if any); and (4) time monitoring intervals (if any).

6.20.4 Choosing CC-Mode for Bright Source Observation

In the past, the continuous clocking (CC) mode (see Section 6.12.3) has been used to mitigate pileup in very bright sources (see Section 6.15.2). ACIS has offered two standard telemetry formats for observations performed in CC mode, faint (F) and graded (G) (see Section 6.14.2). While the CC-F mode choice has been used on occasion, the CC-G mode has been in the program for over 2 Msec of Chandra observing time, primarily to accommodate HETG spectra of bright X-ray binaries. The objective was to use CC mode in order to effectively mitigate pileup and conserve discrete structures such as emission and absorption lines, and edges in the dispersed spectrum. G mode also reduces the possibility of telemetry saturation.

Spectra in CC-mode, however, carry a number of currently uncalibrated artifacts ranging from a wrong spectral shape at wavelength larger than 10 Å, a smeared Si K edge, and charge losses below 3 Å. For bright sources up to about 300 mCrab we recommend the use of TE faint mode with four CCDs S1 - S4 and a 350 row subarray, at the expense of the spectral range above 12 Å ( < 1 keV), some pileup ( < 10%), and some moderate amount of frame drops. The latter can be mitigated by using TE graded mode in extreme cases.

We now have developed methods to calibrate modes similar to, but not identical to, the G mode. Starting with Cycle 12, the G mode has been altered to include some of the flight grades that were previously rejected on board. ACIS now telemeters all flight grades except 24, 107, 127, 214, 223, 248, 251, and 255.

Here are a few guidelines for the use of CC-mode and the HETG. Please note, that in any case spectra above above 12 Å ( < 1 keV) are always affected by some enhanced background:

•: CC-faint mode can be used for sources of about 60 mCrab with no additional measures to restrict event selection
•: CC-faint mode can be used for sources of about 90 mCrab and a zero order exclusion window selection of only every 10th (or less) photon
•: CC-faint mode can be used for sources of about 105 mCrab and a zero order exclusion window selection of only every 10th (or less) photon and an event energy filter starting at 0.3 keV and a range of 10 keV (or less).

For even brighter sources and in the case of HETG spectra, other mitigating actions can be taken, such as shutting off CCDs or moving half the source and dispersion off the array (i.e. using Z-SIM = -11mm). In these cases sources of over 300 mCrab can be accommodated.

For sources significantly exceeding such flux levels, CC-graded mode may be used. Like in the case of CC-faint mode, ASCA grade 7 events are no longer entirely rejected and similarly additional flight grades enter the telemetry stream causing frame drops to increase by about 10% compared to the previous definition. Mitigating strategies similar to those described above may be used.

INDEX

Last modified:12/14/11

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Chapter 6 ACIS: Advanced CCD Imaging Spectrometer