The XRISM Proposers’ Observatory Guide

Gate-Valve Closed Version: February, 2024

Announcement of Opportunity #1
Institute of Space and Astronautical Science (ISAS/JAXA),
X-ray Astrophysics Laboratory
NASA’s Goddard Space Flight Center (GSFC)
and theEuropean Space Agency (ESA)

Copies of this guide are available in html and pdf formats.

Contents
1 Introduction
1.1

2

Purpose and Scope of this Document . . . . . . . . . . . . . . . . . . . . . .

2 Mission Description

3
5

2.1

Spacecraft: Orbit and Attitude . . . . . . . . . . . . . . . . . . . . . . . . .

6

2.2

Overview of XRISM Instruments . . . . . . . . . . . . . . . . . . . . . . . .

6

2.2.1

X-ray Mirror Assembly (XMA) . . . . . . . . . . . . . . . . . . . . .

7

2.2.2

Resolve . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

7

2.2.3

Xtend . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

93 Observation Policies

11

3.1

Data Rights . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

11

3.2

Target Selection and Prioritization . . . . . . . . . . . . . . . . . . . . . . .

11

3.3

Target of Opportunity (TOO) Proposals . . . . . . . . . . . . . . . . . . . .

12

3.4

Pointing Constraints . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

12

3.5

Events After Submission . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

14

4 X-Ray Mirror Assembly (XMA)
4.1

4.2

16

XMA Components . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

18

4.1.1

Mirror Part . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

18

4.1.2

Pre-Collimator . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

18

4.1.3

Thermal shields and heaters . . . . . . . . . . . . . . . . . . . . . . .

21

Expected Performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

21

4.2.1

21

Focal Positions and Optical Axis . . . . . . . . . . . . . . . . . . . .
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iii

4.2.2

Effective Areas . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

22

4.2.3

Vignetting . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

24

4.2.4

Angular Resolution . . . . . . . . . . . . . . . . . . . . . . . . . . . .

24

4.2.5

Stray Light . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

26

5 Resolve

31

5.1

Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

31

5.2

The microcalorimeter. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

34

5.2.1

Photon detection . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

35

5.2.2

Anti-coincidence detector and particle background . . . . . . . . . .

35

Event processing system . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

37

5.3.1

Event triggering . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

37

5.3.2

Pulse height determination . . . . . . . . . . . . . . . . . . . . . . .

38

5.3.3

Event grading: low, mid, and high-res grades . . . . . . . . . . . . .

38

5.3.4

Event grade branching ratios . . . . . . . . . . . . . . . . . . . . . .

40

5.3.5

Pulse Shape Processing limit . . . . . . . . . . . . . . . . . . . . . .

41

5.3.6

Cross-talk . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

41

5.4

Detector gain . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

42

5.5

Line spread function . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

42

5.6

Calibration sources and MXS operations . . . . . . . . . . . . . . . . . . . .

42

5.7

Background . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

43

5.8

Planning observations with Resolve . . . . . . . . . . . . . . . . . . . . . . .

44

5.8.1

Filter selection . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

44

5.8.2

Observing faint to moderately bright point sources . . . . . . . . . .

46

5.8.3

Observing extended sources . . . . . . . . . . . . . . . . . . . . . . .

46

5.8.4

Observing bright sources . . . . . . . . . . . . . . . . . . . . . . . . .

46

5.3

6 Xtend/SXI

47

6.1

Xtend Basics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

47

6.2

Observation Options for General Users . . . . . . . . . . . . . . . . . . . . .

52

6.3

Things to Be Considered for Bright Source Observations . . . . . . . . . . .

53

6.3.1

53

Photon Pileup . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

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6.4

iv

6.3.2

Out-of-Time Events . . . . . . . . . . . . . . . . . . . . . . . . . . .

54

6.3.3

Sacrificial Charge . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

54

SXI in Depth . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

54

6.4.1

X-ray Detection Mechanics and Operation . . . . . . . . . . . . . . .

55

6.4.2

Charge Transfer Inefficiency and Charge Injection Operation . . . .

57

6.4.3

Hot Pixels . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

59

6.4.4

Area Discrimination . . . . . . . . . . . . . . . . . . . . . . . . . . .

59

6.4.5

Supporting Components . . . . . . . . . . . . . . . . . . . . . . . . .

59

6.4.6

Change from the Hitomi SXI . . . . . . . . . . . . . . . . . . . . . .

61

7 Observations of Extended Sources

62

7.1

Challenges of analyzing extended sources. . . . . . . . . . . . . . . . . . .

62

7.2

Facing the challenge: methods and prospects . . . . . . . . . . . . . . . . .

66

7.3

Consequences and general considerations . . . . . . . . . . . . . . . . . . . .

69

8 Resolve Observations of Bright Sources

73

8.1

Challenge: High-Resolution Event Rates . . . . . . . . . . . . . . . . . . . .

74

8.2

Challenge: Event Loss . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

74

8.3

Challenge: Minimizing Cross-talk . . . . . . . . . . . . . . . . . . . . . . . .

75

8.4

Observing Strategies: Filters and Off-Axis Pointing . . . . . . . . . . . . . .

76

8.5

Filtering Cross-talk on the Ground . . . . . . . . . . . . . . . . . . . . . . .

79

8.6

Summary of Observing Recommendations . . . . . . . . . . . . . . . . . . .

82

List of Figures
2.1

The 96 minute XRISM orbit. . . . . . . . . . . . . . . . . . . . . . . . . . .

5

2.2

A schematic of the XRISM spacecraft. Resolve dewar and its XMA are
highlighted in blue, while the Xtend instrument and its XMA are highlighted
in pink. Image credit: JAXA. . . . . . . . . . . . . . . . . . . . . . . . . . .

7

Schematic layout of the two scientific instruments on-board XRISM , Resolve
and Xtend. This is the view looking toward the sky from the back of the
focal plane. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

9

4.1

Image of the XMA with all its components. . . . . . . . . . . . . . . . . . .

16

4.2

Picture of the X-ray Mirror Assemblies, without thermal shields (right:
Resolve-XMA, left: Xtend-XMA). . . . . . . . . . . . . . . . . . . . . . . .

17

Major reflection paths occurring in the XMA structure. (a) Normal double
reflection of the X-ray arriving from the on-axis direction. Incident X-rays
are bent by the angle 4τ in total, and converge to the on-axis focus: (b)Secondary reflection, which arrives at the focal plane only if the incident
angle of the X-ray, θ, measured from the optical axis is in the range τ < θ <2τ : (c) Stray light path that gives rise to the brightest ghost among various
backside reflections, the reflection at the backside of the primary followed by
the normal double reflection. This pattern occurs when the X-ray incident
angle is in the range 2τ < θ < 3τ . From Mori et al. (2005). . . . . . . . . .

19

4.4

Image of a Pre-collimator quadrant used for XRISM XMA. . . . . . . . . .

20

4.5

Image of a thermal shield quadrant used for XRISM XMA. . . . . . . . . .

20

4.6

Cartoon view of the actual configuration of a mirror compared to the ideal
geometry (shown on the left). The middle picture illustrates the shift between the focal point and the detector aim point and the right picture showsin addition the effect of a tilt of the optical axis with respect to the nominal
aim-point direction (HITOMI in-flight calibration plan, 2015.) . . . . . . . .

22

2.3

4.3

v

Proposers’ Observatory Guide
4.7

vi

Left: Layout of Xtend FoV (black) superimposed with Resolve FoV (cyan).
This is a look-up view. The expected Xtend aimpoint is located from ∼5.0′(DETX) and ∼ 5.4′ (DETY) from the edges of the CCD. The ground calibration reported that Xtend is on-axis at aim point (less than 3” shift).
Right: The expected aimpoint of Resolve as the center of Resolve FoV. The
ground calibration reported that optical axis of the XMA-Resolve is slightly
shifted. In-flight calibration is on-going. . . . . . . . . . . . . . . . . . . . .

23

Effective area versus energy, for Xtend-XMA (black) and Resolve-XMA (red)plotted on the linear scale (left) and on the log scale (right). Note that thesesimulated curves are based on ground measurements and the XMA-only
effective areas, thus not including QE of the detectors, etc.. . . . . . . . . .

23

Left: Vignetting curves of all energies, for Xtend-XMA, for azimuthal 0-180
deg, fitted with the sum of Lorentzian and Gaussian models to the datameasured on the ground. Right: Look-down illustration of the mirror rollangle (different from the spacecraft roll angle) and the off-axis angle usedduring the ground calibration measurements (see Boissay-Malaquin et al.
(2022)). The coordinates of the satellite and detector are shown, as well as
the boundaries of the XMAs’ quadrants. . . . . . . . . . . . . . . . . . . .

24

4.10 Vignetting curves of all energies, for Resolve-XMA, for roll angle (azimuthal
angle) 0-180 deg, fitted with 1D Lorentzian model fitted to the data measured on the ground. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

25

4.11 Vignetting curves at 6.4 keV fitted to the data points measured on the
ground, for Resolve-XMA (left) and Xtend-XMA (right), for 8 roll (azimuthal) angles (see the right panel of Figure 4.9). . . . . . . . . . . . . . .

25

4.12 Field of View of the XMAs, as a function of energy. The FoV is defined here
as the FWHM value determined by fitting the vignetting curves with a 1D
Lorentzian model, and does not take the physical size of the detectors into
account. The FoV is the average value of the FWHM in all directions. The
data points are from the ground measurements. . . . . . . . . . . . . . . .

26

4.8

4.9

4.13 Focal plane images of Resolve-XMA (top) and Xtend-XMA (bottom) at thesix energies of 1.49, 4.50, 6.40, 8.05, 9.44, and 11.07 keV (Tamura et al.,2022). The image size is 6.7′ × 6.7′ . . . . . . . . . . . . . . . . . . . . . . .
4.14 One-dimensional PSFs of Resolve-XMA (left) and Xtend-XMA (right). The
six PSFs at 1.49 (red), 4.50 (green), 6.40 (blue), 8.05 (light blue), 9.44(violet), and 11.07 keV (dark blue) are plotted (Tamura et al., 2022). Theseplots are based on the ground measurements. . . . . . . . . . . . . . . . . .

28

29

Proposers’ Observatory Guide
4.15 The EEF of the Resolve-XMA (left) and the Xtend-XMA (right). The six
EEFs at 1.49 (red), 4.50 (green), 6.40 (blue), 8.05 (light blue), 9.44 (violet),and 11.07 (dark blue) are plotted (Tamura et al., 2022). These plots are
based on the ground measurements. . . . . . . . . . . . . . . . . . . . . . .
4.16

vii

29

Stray light images of Resolve-XMA (top) and Xtend-XMA (bottom) at 30′
off axis angle (left) and 60′ off axis (right) (Tamura et al., 2022). The X-ray
energy is 1.49 keV. The image size is 17.8′ ×17.8′ . The FOV of Resolve isshown in left top panel. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

30

A schematic of the Resolve microcalorimeter array focal plane. Dimensions
shown are the design values, but most actual absorbers were 0.818 mm wide,
and the gaps were correspondingly smaller on average. . . . . . . . . . . . .

33

Effective area of Resolve, which includes the gate valve transmission (Midooka et al., 2021). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

33

5.3

Transmission curve for the non-selectable five-filter blocking filter stack. . .

34

5.4

The concept of an X-ray calorimeter is shown in the left panel, whereby an Xray deposits energy in an absorber that is measured by a thermometer after
the energy thermalizes. Subsequently, the absorber returns to its quiescenttemperature and the system is ready to detect and measure the next X-ray.
Right panel shows a schematic of the temperature of an absorber, which isa part of Resolve pixels, as a function of time. Upon absorbing the photon,
the temperature goes up proportionally to the energy of the incident photon,and then goes down with time constant proportional to the heat capacity
and thermal link conductivity. . . . . . . . . . . . . . . . . . . . . . . . . .

36

Schematic arrangement of the XRISM Resolve microcalorimeter array andthe anti-coincidence detector (c.f. Porter et al., 2010). The microcalorimeter
array sits on top of the anti-co detector. . . . . . . . . . . . . . . . . . . . .

36

Typical shape of a pulse triggered by an incident X-ray. Left panel showsa single photon, while right panel shows a situation where second photon isdetected before the pixel equilibrates after the first event. The undershootfollowing the pulse happens because the electronics are AC coupled. For
details, see Section 5.3.1. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

37

Comparison of the fluorescent Mn Ka spectra for different event grades fromthe Resolve ground calibration program. . . . . . . . . . . . . . . . . . . .

39

Sketch illustrating the grading of Resolve events. The five cases described in
the text (high primary, medium primary, medium secondary, low primary,and low secondary) are considered. For details see Ishisaki et al. (2018). . .

40

5.1

5.2

5.5

5.6

5.7
5.8

Proposers’ Observatory Guide
5.9

viii

The event grade branching ratio is changed by the incoming rate. At low
rates, most events will achieve calorimeter resolution (Hp or Mp ). For brightsources, resolution is degraded. . . . . . . . . . . . . . . . . . . . . . . . . .

41

5.10 Instrument response to a monochromatic emission line at 5.414 keV, demonstrating the Gaussian core which dominates the line shape and the electron
loss continuum. For more details, see Eckart et al. (2018) . . . . . . . . . .

43

5.11 Spectrum of the NXB background measured in-orbit in the Hitomi SXS
microcalorimeter. For details, see Kilbourne et al. (2018a). . . . . . . . . .

44

5.12 Filter wheel filter configuration arrangement, viewed looking up from thedetector toward the sky. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

45

5.13 Transmission curves for the selectable filters, the Neutral Density (ND) andBe-filter. For filter effective area curves see Figure 8.3. . . . . . . . . . . .

46

6.1

6.2

6.3

X-ray images of the Perseus cluster obtained with the Hitomi observatory(Nakajima et al., 2018). (a) Full X-ray CCD (SXT+SXI) image overlaid
with the micro-calorimeter (SXT+SXS) FOV in cyan. (b) Magnified CCD
image of the central region overlaid with the SXS pixel boundaries in cyan.
(c) Micro calorimeter (SXT+SXS) image of the same region. . . . . . . . .

48

Left — Schematic layout of the SXI CCDs looking toward the sky from theback of the focal plane. The 4 CCD chips are aligned in a 2×2 format. Each
CCD has two electrically separated segments, whose boundary is shownby the dashed line. Each segment has two readout nodes for redundancy(the black dots with the A/B/C/D labels). Each chip has an additionalframe store area on the readout node side. The filled gray half circles show
approximate areas illuminated by the calibration sources. The red box shows
the SXS FOV and the dotted orange circles show the off-axis angles. The
distances from the CCD edges are inflight measurement values. There is atiny area on the far side corner from the aim point shown in yellow, where acamera body structure blocks the view of the sky. Right — SXI CCD sensor
photo (Nakajima et al., 2020, Fig. 4). Gold-coated metal frames cover theframe store areas. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

49

Left — SXI CCD Quantum efficiency measured with the ground experiments(black). The plot also separately shows the quantum efficiency of a CCD
chip combined with the transmission of the Optical Blocking Layer (QE +OBL, red) and the transmission of the Contamination Blocking Filter (CBF,
green). Right — Effective area of the Xtend telescope, including the XMA
component. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

50

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6.4

ix

Left — Hitomi SXI NXB spectrum normalized by the physical pixel area.
The plot also shows a Suzaku XIS-BI NXB spectrum for the same physicalpixel area. Right — Hitomi SXI total (NXB + sky) background spectrum
normalized by the effective area and solid angle of the sky. The plot also
shows the CXB spectrum and background spectra of other X-ray CCD instruments (Nakajima et al., 2018). . . . . . . . . . . . . . . . . . . . . . . .

51

Simulated SXI image of the Perseus cluster with the 1/8 window mode. The
green boxes show the approximate chip boundaries. . . . . . . . . . . . . .

52

6.6

External View of SXI-S . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

56

6.7

Schematic diagram of the charge transfer process on a CCD. (a) At the end of
a frame exposure, X-rays or charged particles produce charge clouds aroundthe reacted pixels (red dots). (b) The CCD clock driver transfers the charge
clouds vertically to the FS region. It also injects artificial charges to the top
serial register every N vertical transfers (N =80 in the prelaunch default
setup), which move down to the IA with the subsequent vertical transfers(red lines). (c) Once all IA charges move to the FS area, the CCD clock
driver only transfers charges in the bottom FS row horizontally for readout
and then moves all FS charges down by one pixel for another horizontal
readout. It repeats this process until all FS charges are read out. IA takes
another exposure during this readout process. (d) Another charge transfer
process begins. The CI spacing and X-ray cloud sizes are not to scale. . . .

57

Results of a ground test of charge injection (CI) operation in September
2021 with an SXI-S CCD engineering model. In this test, 5 keV charges
were injected to the CCD chip by every 80 binned rows. Left — Peak Pulse
Height (PH) values of Mn Kα source events with/without (red/black) CI
with the ACTY row, equivalent to twice the number of parallel transfers.
CI prevents PH (i.e. charge) reduction. It produces a saw-tooth pattern
because it works better on pixels following near a CI row. Right — 55 Fe
radioactive source spectra accumulated from all events on the CCD chip.
CI recovers the line energy and the energy resolution. . . . . . . . . . . . .

58

55Fe calibration source spectrum, obtained during the thermal vacuum test
in August 2022. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

60

Shape of the PSF and its effective impact on Resolve. Top panels: case of apoint-like source placed at the center of the field of view (left), translating
into significant internal SSM on the Resolve pixels (right). Bottom panels:
case of a point-like source placed directly outside of the field of view (left),translating into significant external SSM on the Resolve pixels (right). “BLCCD” stands for “Beam-Line CCD”. . . . . . . . . . . . . . . . . . . . . . .

64

6.5

6.8

6.9

7.1

Proposers’ Observatory Guide
7.2

7.3

7.4

8.1

x

Simplistic case of spatial mixing for a circular extended source whose left andright halfs are redshifted and blueshifted, respectively (uppermost panel).
The two intermediate panels show, from that region (grey dotted circle),a distribution of all the incoming (redshifted or blueshifted) photons eventually reaching the Resolve detector (with yellow square and dashed lines
lines indicating the Resolve detector boundaries).The XRISM PSF effectively mixes photons from each part into the other, resulting in 15-17% ofcross-contamination (two bottom panels, showing the fractional number ofphotons). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

65

Spectral mixing effects on the Fe-K line at 6.4 keV as seen by Resolve,
assuming the simplistic case presented in Figure 7.2 (V1 = -235 km s−1 ;V2 = +235 km s−1 ). The top four panels assume a line width of σ = 1 eV
(FWHM = 110 km s−1 ), separating cases for contamination of equal fluxes(left panels) vs. cross-contamination of 16% representing the regions shownin Figure 7.2 (right panels), as well as theoretical models (upper panels) vs.
more realistic profiles convolved by the Resolve RMF (lower panels). The
bottom four panels show the same effect for a larger line dispersion width(σ = 4 eV; FWHM = 440 km s−1 ). . . . . . . . . . . . . . . . . . . . . . .

71

Illustration of the tasks performed by the routine xaarfgen—and the utilityof the latter in computing SSM coefficients for the analysis of extended
sources. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

72

Electrical cross-talk “child” pulses occur in electrically adjacent pixels, so thecross-talk distribution does not follow the shape of the PSF. Left: Resolve
pixel map showing electrically neighboring pixels. A parent pulse in pixel
18 (blue) will produce a child pulse in pixel 19 (light blue), but not in pixel
17 because it is in a different and electrically isolated quadrant. Likewise,a parent pulse in pixel 31 (orange) will produce child pulses in both pixels
30 and 32. Right: Cross-talk pulses (blue) have a shape based on the time
derivative of the parent pulse (red) and a much smaller signal, with a peak
amplitude of 0.6% of the pulse height. The child pulses are superimposedon the data stream in their pixels; when a child pulse in pixel 19, induced
by a parent pulse in pixel 18, occurs around the same time as a parent pulse
in pixel 19, the pulse shape is distorted. This leads to an erroneous energy
measurement. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

77

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8.2

8.3

8.4

xi

Left: Increasing contamination by cross-talk child pulses both broadens lines
and shifts their centroids. At high count rates, this can significantly alter
the energy resolution in individual pixels, as shown here for an example FeKα line. Right: As an example, this Resolve pixel map shows the energy
resolution (i.e., broadening) in each pixel when starting from ∆E = 5 eV andobserving a source four times brighter than the Crab Nebula at the array
center. The four exterior pixels with the worst resolution are electricallyadjacent to the central four pixels. Note that the map is not weighted by
the event rate or grade: despite the better energy resolution in the central
pixels, there are hardly any Hp+Mp events from those pixels. . . . . . . . .

78

Resolve effective area curves with the gate valve closed for no filter (black),
the Be filter (red), and the neutral-density filter (blue). The neutral-density
filter reduces the count rate by about a factor of four across the bandpass.
For filter transmission curves see Figure 5.13. The Be filter is optimized forsuppressing (continuum) events at low energy in favor of preserving highresolution events in the Fe-K complex. . . . . . . . . . . . . . . . . . . . . .

80

These six plots show a mock image (top left), raw incident count rate (topcenter), processed Hp+Mp rate (top right), cross-talk child pulse rate (bottom left), effective spectral resolution (bottom center), and centroid shift(bottom right) in each pixel for a 10−7 erg s−1 cm−2 source observed nofilters. The top right quadrant is sacrificed to maximize the Hp+Mp rate in
the other three quadrants, as each quadrant is processed separately in thePSP. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

81

List of Tables
2.1

XRISM capabilities at a glance. . . . . . . . . . . . . . . . . . . . . . . . . .

10

4.1

XMA parameters and their characteristics. . . . . . . . . . . . . . . . . . . .

17

4.2

The main characteristics of the XMA pre-collimator. . . . . . . . . . . . . .

20

4.3

Effective area (in cm2 ) for both XMAs (with 68% statistical errors),atdifferent energies E, for the whole telescope axis. From Boissay-Malaquin etal. (2022) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

22

4.4

The HPD obtained at seven energies. From Tamura et al. (2022) . . . . . .

26

4.5

The effective area at 1.49 keV for off-axis sources due to stray light. From
Tamura et al. (2022) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

27

Summary of the expected Resolve detector properties based on the groundcalibration. HPD stands for Half Power Diameter. . . . . . . . . . . . . . .

32

Definition of event grades. Calorimeter-grade energy resolution is obtainable
only for events labeled Hp and Mp, marked in bold font . . . . . . . . . . .

40

6.1

Xtend Characteristics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

48

6.2

Available Window/Burst Mode Options . . . . . . . . . . . . . . . . . . . .

52

6.3

SXI Characteristics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

55

8.1

Resolve Count Rates . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

79

8.2

Bright Source Observing Recommendations. . . . . . . . . . . . . . . . . . .

82

5.1
5.2

1

Chapter 1

Introduction
The X-Ray Imaging and Spectroscopy Mission (XRISM ) is a mission led by Institute ofSpace and Astronautical Science (ISAS) of Japan Aerospace Exploration Agency (JAXA)
in collaboration with the National Aeronautics and Space Administration (NASA), theEuropean Space Agency (ESA), and other partners in Japan and around the world. Its
main purpose is to perform high spectral resolution, high throughput imaging X-ray spectroscopy of celestial X-ray sources. XRISM was lanched on 2023 September 6 (UT) on
a JAXA HII-A rocket. As of this writing, the initial in-orbit check-out and commissioning phase of XRISM has just concluded, and 6 months of Performance Verification (PV)observations, whose targets have been chosen by the XRISM Science Team, have started.
Subsequently, XRISM will become a general purpose observatory open to all astronomers,
with annual calls for proposals issued in parallel by JAXA, NASA, and ESA. In addition,
all PV data will enter public archive 1 year after the conclusion of the PV phase; all subsequent observational data will enter public archive after an initial 1-year exclusive use
period.
The scientific payload of XRISM consists of two instruments, Resolve and Xtend.
The Resolve instrument has a Gate Valve (X-ray aperture door) that protects thedetector on the ground as well as during launch and early operations, and is meant to beopened during commissioning. However, the Gate Valve has not opened, thereby shifting
the nominal energy band of Resolve from 0.3–12 keV to 1.7–12 keV and lowering theeffective area. While the XRISM team will continue to assess different approaches to
opening the Gate Valve, the Cycle 1 program will be carried out with the closed Gate
Valve configuration, and therefore this document describes Resolve in this configuration.
Note that the Resolve Gate Valve does not affect the Xtend instrument.
• Resolve, a soft X-ray spectrometer, which combines a lightweight X-ray Mirror Assembly (XMA) paired with an X-ray microcalorimeter spectrometer, and provides
non-dispersive ∼5 eV energy resolution in the 1.7–12 keV bandpass with a field of2

Proposers’ Observatory Guide3view of ∼3 arcmin.
• Xtend, a soft X-ray imager, consisting of an array of four CCD detectors that extend
the field of the observatory to 38×38 arcmin2 over the energy range 0.4–13 keV.
Xtend uses a lightweight XMA identical in design to that used for Resolve.
These instruments are nearly identical to the two soft X-ray instruments flown on
Hitomi, a mission that was lost shortly after launch in 20161 . The Soft X-ray Spectrometer(SXS) on Hitomi and Resolve on XRISM are sufficiently similar, as are the Soft X-ray
Imager (SXI) on-board Hitomi and Xtend on XRISM , that the papers based on Hitomi
data serve as an excellent guide in preparing XRISM observing proposals. Note also thatall Hitomi observations were carried out with the SXS Gate Valve closed. In addition, data
formats as well as data analysis and simulation software for XRISM are based largely on
those developed for Hitomi.

1.1

Purpose and Scope of this Document

This document is intended for researchers who have some general understanding of spacebased astronomy and X-ray observations. It should provide the readers with a general
understanding of the characteristics of Resolve and Xtend, including on-board data handling, calibration, and operational limitations. This document also explains the basics of
how XRISM will be operated, and how the observing time will be allocated. Based on
this document, the reader should be able to determine whether XRISM is the appropriate
instrument to address their scientific question, and how to design appropriate observing
program for the proposal.
This document is not intended as a data analysis guide. Neither is it intended toprovide in-depth descriptions of the instruments. In addition, the mechanical details of
proposal submission will be explained elsewhere.
The document is organized as follows. Chapter 2 summarizes the principal characteristics of the detectors on-board XRISM . Chapter 3 covers how the observing time will beallocated. This includes data rights and time available for General Observers (GOs), as
well as policies regarding Targets of Opportunity (TOOs). In addition, observational constraints due to the orbit, Sun angle, and the pointing accuracy are described. The proposal
process is reviewed, including how to submit a proposal, how they will be evaluated, andhow observations will be scheduled, performed, and the results disseminated.
The subsequent three chapters describe XMA (Chapter 4), Resolve (Chapter 5), and
Xtend (Chapter 6). Finally, Chapter 7 and Chapter 8 focus on particular caveats associated1

Note that XRISM does not carry the equivalents of the hard X-ray and soft gamma-ray instruments
on-board Hitomi.

Proposers’ Observatory Guide

4

with, respectively, extended sources and bright sources, which may thus be of importancefor many proposers.
Disclaimer:
The current version of this document was prepared largely based on the best prelaunch understanding of the observatory and its instruments, with some early indications
from early commissioning activities. Our understanding of the performance of the mission
is expected to evolve, quantitatively if not qualitatively. Any significant new information
will be made available at our websites and through other communication channels.

Chapter 2

Mission Description
This chapter is a brief introduction to the satellite and its instruments and is intended
to provide a quick overview. Reading it thoroughly should provide the reader with thenecessary information to understand the capabilities of the instruments at a level sufficient
to judge if XRISM is suitable for a scientific project, and to prepare the feasibility sectionof a XRISM proposal.

Figure 2.1: The 96 minute XRISM orbit.

5

Proposers’ Observatory Guide

2.1

6

Spacecraft: Orbit and AttitudeXRISM is in many ways similar to previous Japanese X-ray astronomy satellites (e.g.,
Suzaku) in terms of orbit, pointing, and tracking capabilities. XRISM is placed in a nearcircular orbit with an apogee of ∼575 km, an inclination of ∼31 degrees, and an orbital
period of about 96 minutes (Figure 2.1). From this altitude, most targets are occulted
by the body of the Earth for about 1/3 of each orbit, except those near the orbital poles.
Additionally, during passages through the South Atlantic Anomaly (SAA), the instruments
will not produce usable data. Thus a 1 day observation of a target typically results in about
40 ks of good on-source time.
The maximum slew rate of the spacecraft is 4.5 degrees/min, and settling to the final
attitude takes ∼ 10 minutes, using the star trackers. Given the relatively long slew time,
the normal mode of operation is to point at a single target for at least 1/4 day (10 ks net
exposure time) and often several days, without interleaving multiple targets.
Objects near the spacecraft orbital poles will not suffer Earth occultations (i.e., they arein the continuous viewing zone, CVZ), thus can be observed at a much higher efficiency.
The orbit is predicted to precess with a period of about 2 months. The orbital poles
will be at declinations of ±59 (90 − 31) degrees and a right ascension that changes withthe precession phase. Because the evolution of XRISM orbit is influenced in part by theatmospheric drag, accurate long-term prediction for the times when a given target is in
CVZ is not feasible.
The XRISM instruments are oriented perpendicular to the solar panels, which are fixed
to the spacecraft, and must remain pointed to within 30 degrees of the Sun. Therefore,targets can be observed only when their Sun angle is in the 60–120 degree range. This
also implies a constraint on the roll angle: when the target is at a Sun angle close to 90
degrees, there is a ±30 degree range of possible roll angles. When observed close to thelimits of Sun angles, the roll angle is tightly constrained.
The overall observing efficiency of the satellite is expected to be about 45%. Proposals
should specify the desired good on-source exposure but without accounting for instrumental
dead times due to high count rates, if relevant.

2.2

Overview of XRISM Instruments

XRISM carries two instruments, Resolve and Xtend, that are co-aligned and operate simultaneously, each behind its own X-ray Mirror Assembly (XMA). Here we provide abrief overview, leaving further details to chapters dedicated to these in turn. We show aschematic of the XRISM spacecraft in Figure 2.2.

Proposers’ Observatory Guide

7

Figure 2.2: A schematic of the XRISM spacecraft. Resolve dewar and its XMA are highlighted in blue,
while the Xtend instrument and its XMA are highlighted in pink. Image credit: JAXA.

2.2.1

X-ray Mirror Assembly (XMA)

The two near-identical X-ray Mirror Assemblies (XMAs) are used to focus X-ray photons
on the two detectors. Both XMAs are composed of 203 thin reflector shells that use
conical approximations of Wolter-I type optics, thereby resulting in a light-weight, large
collecting area (∼580 cm−2 at 1.5 keV and ∼420 cm−2 at 6.4 keV) X-ray telescopes.
In terms of imaging quality, XMAs represent an incremental improvement over previous
telescopes of the same type, such as those used for Suzaku. Half-power diameter (HPD)
of 1.3′ is measured on ground for Resolve-XMA, and 1.5′ for Xtend-XMA. Stray light, or
contaminating X-rays from outside of the field of view (FOV), is reduced by a stray-lightbaffle also called a Pre-Collimator (PC) placed above each reflector. A Thermal Shield(TS) is attached in front of the pre-collimator to stabilize the thermal environment of thetelescope.
A more detailed descriptions of XMA can be found in Chapter 4.

2.2.2

Resolve

Resolve consists of an X-ray micro-calorimeter array of 6x6 pixels placed at the focus of an
XMA, covering a ∼ 3′ × 3′ arcmin2 region of the sky. It is almost identical to the Hitomi
Soft X-ray Spectrometer. Its unprecedented non-dispersive high spectral resolution (∆Earound 5 eV) capability is achieved by cooling the detector to 50 milli-Kelvin and measuringProposers’ Observatory Guide

8

the temperature rise when an X-ray photon is absorbed by the detector element.
The ability of the Resolve instrument to determine the photon energy depends onrecording the complete history of temperature rise and decay associated with each photon,
isolated from all other events, to be compared with templates. This leads to the conceptof event grade. When the count rate is high, the fraction of high resolution events will
decrease. Moreover, for very bright sources, there is a limit imposed by the Pulse Shape
Processor (PSP) which is unable to process events at a rate above ∼200 per second (seeChapter 8 for more details). For both these reasons, it may be necessary and/or advisable
to reduce the count rate of the brightest X-ray sources using a filter or offset pointing.
For the former, both a neutral density and a (soft X-ray absorbing) beryllium filters areprovided.
The effective area is ∼180 cm2 at the 6–7 keV K-band, considerably larger than any
other high-resolution X-ray spectrometers. The imaging capability is limited both by the
pixel size and the imaging performance of the XMA. The calorimeter array is backed
by an anti-coincidence detector that enables rejection of background events. The milliKelvin temperature in the Dewar is achieved by a cooling system that includes adiabatic
demagnetization refrigerators, He coolants, and mechanical coolers. The cooling systemallows redundancy and even cryogenic-free operation. For calibration, 55 Fe sources (onecontinuously illuminating a dedicated calibration pixel, and one on the filter wheel thatcan be inserted into the light path) and Modulated X-ray Sources (MXS) are available.
The best strategy to utilize these calibration sources, as well as the in-orbit performanceof Resolve in terms of gain and line spread function, are currently being investigated, but
gain calibration accuracy of better than ∼1 eV is expected.
In comparison with grating instruments, the two key advantages of Resolve are:• High spectral resolution for spatially extended sources: the energy resolution of Resolve is unchanged regardless of spatial extent, whereas those of grating instruments
become degraded.
• High spectral resolution in its bandpass, particularly the Fe K region. Note that theenergy resolution of Resolve is a slowly rising function of photon energy and is better
than ∼5 eV for much of its bandpass except well above the Fe K band, resulting
in a higher resolving power at higher energies. The opposite is the case for grating
instruments whose wavelength resolution is roughly constant for a given order of agiven instrument.
A more detailed descriptions of Resolve can be found in Chapter 5.

Proposers’ Observatory Guide

2.2.3

9

Xtend

Xtend is an array of four X-ray CCDs behind the other XMA. It has an imagingspectroscopic capability of a wide field (38×38 arcmin2 ) and a medium energy resolution
in the 0.4–13 keV band (comparable to Suzaku XIS, Chandra ACIS, and XMM-Newton
EPIC). There is a gap of ∼20 arcsecond between the chips (see Figure 2.3 for the layout).
The nominal pointing position is placed at the center of Resolve FoV, which is offset fromthe center of Xtend by ∼ 5′ .
~40’’ (23 pix)

18.9’ (640 pix)D

CCD_ID=0

CCD_ID=2(CCD1)(CCD3)

B

B

C
15

is

-ax

10

xis

ﬀ-a

’o

D

D

5’

is

-ax

oﬀ

A

5.4’

~40’’ (23 pix)

A

ﬀ
’o

18.9’ (640 pix)

C

A

5.0’
C

Resolve FoV

B(3.05’ x 3.05’)

B

C

DETY

SAT X

DETY
Roll

A

DETX

CCD_ID=1

CCD_ID=3(CCD2)(CCD4)SAT Y
D

55Fe calibration sources

Figure 2.3: Schematic layout of the two scientific instruments on-board XRISM , Resolve and Xtend. Thisis the view looking toward the sky from the back of the focal plane.

The background level of Xtend is expected to be low and stable benefiting from thelow-Earth orbit of the satellite. Above ∼7 keV, the use of a thick depletion layer should
lead to a significant background reduction compared to the Back-Illuminated (BI) CCD
in Suzaku XIS. All four CCDs on-board XRISM are BI devices and are more resistant tomicro-meteorite impacts than Front Illuminated (FI) devices .
A more detailed descriptions of Xtend can be found in Chapter 6.

DETX

Proposers’ Observatory Guide

10

Table 2.1: XRISM capabilities at a glance.

S/C

XMA

Orbital apogee
Orbital period
Observing efficiency
Energy Range
Effective Area
Angular Resolution

Resolve

Xtend

Energy Range
Efective Area
Energy Resolution
FOV
PixelsFilters
Background
Energy Range
Effective Area
Energy Resolution
FOV

575 km
96 minutes∼ 45%
0.3–15 keV∼580 cm−2 at 1.5 keV∼420 cm−2 at 6.4 keV
1.3′ HPD for Resolve-XMA
1.5′ HPD for Xtend-XMA
1.7–12 keV∼180 cm2 at 6 keV.
∼5 eV FWHM
3.05′ × 3.05′ arcmin2
6×6 array
Open, ND, and Be
0.8 × 10−3 ct s−1 keV−1
0.4–13 keV∼350 cm−2 at 1.5 keV∼300 cm−2 at 6.0 keV∼180 eV FWHM at 6 keV
38′ × 38′ arcmin2

Chapter 3

Observation Policies
Following the PV phase, XRISM will enter the General Observer (GO) phase, which it
will remain in for the lifetime of the mission, with annual Announcements of Opportunity(AO) released simultaneously by JAXA, NASA, and ESA. The institutional affiliationof the proposer determines which agency they should propose to. Persons at Japaneseinstitutions should submit to the JAXA solicitation; those at US or Canadian institutions
should submit to the NASA solicitation; and those with affiliations in ESA member andcooperating states should submit to the ESA solicitation. Persons with affiliations outside
Japan, the US, Canada, or an ESA member state should submit to the JAXA solicitation,
which reserves some time for such international proposals.

3.1

Data Rights

The data rights policy for XRISM is similar to previous missions. The normal exclusive
use period for GO data is one year. This period may be waived if desired by proposers,
but may not be extended.

3.2

Target Selection and Prioritization

The details of the AO are described more fully in the Call for Proposals (CfP). As withprevious missions, proposals will be received by the relevant agency and evaluated bysubject matter experts organized into review panels that make up the Time Allocation
Committee.
Accepted targets will be classified into three categories: A, B, or C. Priority A andB targets are guaranteed to be observed, with A targets receiving a higher priority forobservation scheduling. Priority C targets will be used as fillers when there are gaps in11

Proposers’ Observatory Guide12

the schedule. Priority A and B targets that have not been completed by the end of anobserving cycle are carried over into the next cycle. Priority C observations that are notobserved during that observing cycle are not carried over. The goal of the target selection
process is to produce a combined list of targets representing approximately 140 percentof the available observing time, split among A (50 percent), B (40 percent) and C (50percent) priority observations.
During Cycle 1, category A and B targets will be considered complete if 90% (forA targets) or 70% (for B targets) of the proposed time is obtained on the source. In
general, supplementary observations will be performed for A or B targets that do not meetthe completion criteria, although it may not always be possible in case of time criticalobservations.
During the GO phase, data from calibration and TOOs requested outside the proposal
process (see below) will not be considered proprietary and will not have an exclusive use
period.

3.3

Target of Opportunity (TOO) Proposals

TOO proposals are allowed for XRISM through the normal proposal process, although
they must be ranked as Priority A to be accepted. Proposals with TOOs should not
be mixed with non-TOO targets. TOO proposals are allowed for short-lived events in
known objects, the timing of which is uncertain. The trigger criteria must be explicit andquantifiable, and stated in detail in the proposal text; a brief summary should appear in
the “Remarks” section of the target form. In addition, TOO proposals must provide anestimated probability of a successful trigger during the AO period. It is the PI’s
responsibility to notify the XRISM project when the criteria are met. Generic TOOs
without a specific target (such as “a nearby supernova”) will not be accepted
in Cycle 1. The number of targets in TOO proposals should not exceed 10.
These policies are subject to change in future cycles.

3.4

Pointing Constraints

The Solar panels on the XRISM satellite are fixed. This places a restriction on the pointing
direction with respect to the satellite-Sun line: the Sun angle constraint during Cycle 1 is
60–120 degrees. This means that at any given time of the year, only a swath of the sky
60 degrees wide is accessible for astronomical observations, and thus most celestial sources
are available for observations for about 60 days every 6 months. If a specific observing
date or a coordinated observation with other missions is required, the proposer must first
determine if the observation is possible. This can be done using the “Viewing” tool on the
HEASARC website:Proposers’ Observatory Guide

13

https://heasarc.gsfc.nasa.gov/cgi-bin/Tools/viewing/viewing.plLong (>1 day) observations are the norm for XRISM . A large number of short observations is an inefficient use of the satellite because of the unusable time during slews andattitude settling. The pointing is expected to be accurate to 0.3′ and can be reconstructed
to better than 0.2′ , except during the initial settling period of up to 10 minutes. Moreover,there is a limit on the number of slews that can be uploaded to XRISM . For these reasons,a minimum exposure time of 10 ks has been set for all proposed observations. For Cycle 1,
observations will be limited to 300 ks per pointing with the total not to exceed 600 ks per
proposal.
Contiguous observations, i.e., observations not disrupted by the observation of anothertarget, are generally guaranteed up to exposures of not more than 100 ks. This limitationis due to moon light constraints for the star trackers’ field of view, conflicts with othertime critical observations, and other operational/planning difficulties. While the operationteam does accept requests for uninterrupted observations longer than the 100 ks, these areconducted on a best-effort basis.
Even during shorter pointed observations, there will typically be interruptions dueto the location of XRISM in a low Earth orbit: Normally, a target will be occulted by
the Earth for ∼30 minutes every satellite orbit. In addition, XRISM will pass throughthe South Atlantic Anomaly (SAA) during parts of 5 or 6 orbits every day. Due to theharsh radiation environment of the SAA, scientific observation is not possible during SAA
passages. There are other variations in the particle background, depending primarily on
the Cut-Off Rigidity (COR)1 . Please check the XRISM website for the most up-to-datesuggestions to reduce times of high background while maximizing the science.
There are also orbital constraints on the orientation of the projected Resolve and Xtend
FOV on the sky. Since both are approximately square, specific roll angle is not expected to
be a common requirement. However, if a specific roll angle is scientifically advantageous,the proposer must first determine if it is allowed. Interested proposers are directed to the
XRISM GOF web site for details. Then the required roll range can be entered on the
RPS proposal form. For objects close to the ecliptic poles it is possible to arrange for any
roll angle by scheduling observations at a specific time, but for those located close to theecliptic, Resolve and Xtend will project on the sky in a nearly fixed orientation. Notethat any roll constraint will make a proposal time critical.
It is possible to specify the time of observations (time critical, or TC, observations)in order to observe during specific phases or for simultaneous observations. Monitoring
observations (repeated observations with a specified interval) or roll-angle constrained observations are also considered time critical, and must be so flagged on the proposal form.
The total accepted time of TC and TOO proposals will be restricted; the numerical limit
1

XRISM is protected from solar and cosmic-ray particles by the geomagnetic field. The COR is anindicator of the minimum momentum required for a particle to reach a specific location, given the averagegeomagnetic field configuration.

Proposers’ Observatory Guide

14

will be evaluated based on experiences but may be of order 15% of the total observation,and proposers should justify their requirements carefully.
Overall the proposers are strongly urged to provide the most accurate information
possible. All information that is indispensable for operation planning should be provided
on the electronic forms. The PIs are advised to utilize the “Remarks” area if they havedetailed requests which cannot be expressed with the check boxes/pull-down menus.

3.5

Events After Submission

After the XRISM proposal deadline, there will be three independent proposal reviews for
the NASA, JAXA, and ESA proposal solicitations. Each review will create a target list
from the proposed observations, ranking the accepted targets as category A, B, or C. Only
category A and B targets are guaranteed to be observed. As stated above, TOOs and time
critical observations are only accepted within some limit (approximately 15% of the totaltime) – however, the project will review this limit before the final target selection is made.
Category C targets will be observed as time permits. An international merging committee
will collate the three target lists and produce a single, unified list. Overlaps among targets
accepted by multiple agencies will be resolved, either by merging the investigations (ifboth parties are willing) or by choosing one. In the latter process, the priority given by thenational reviews, as well as the lengths of the accepted observations, will be considered.
The final target list will be ∼ 40% oversubscribed. Category A targets will have 50% of
the available time, category B 40%, and category C 50%.
The Cycle 1 allocation of GO time among the mission partners is 48% for Japanese
investigations (including 4% for international proposals), 44% for US investigations (including Canada), and 8% for ESA investigations.
Even though observations are scheduled to acquire roughly the approved exposure
time and although this will usually be achieved with XRISM , occasional losses of usable
observation time are inevitable. As mentioned above, Category A observations will bedeemed complete when they have received at least 90% of the approved time. Note that
this will be judged based on the good time intervals of the cleaned Resolve event files
after the standard screening. Additional observations will be scheduled automatically for
those non-time-critical targets the observations of which are considered incomplete by the
project scientist at ISAS. In the case of time-critical observations which are incomplete or
unusable, it will be the PI’s responsibility to determine the best course of action.
Each PI will be assigned a contact scientist, either at ISAS or the NASA XRISM GOF,
who will work with the PI to assure the maximum science return. This will include doublechecking coordinates, count rates and finalizing configurations (nominal pointing, modes,etc.). It is important to note that once an observation has been scheduled, any delay inresponding to questions from the contact scientist may result in targets being removed

Proposers’ Observatory Guide

15

from the schedule.
Once the observation has been completed, the data will be promptly run through theprocessing pipeline and put into both the US and Japanese archives in encrypted form.
The PI will be sent the decryption key, along with instructions on how to download anddecrypt the data.
With the exception of the code that converts raw binary telemetry into FITS format
files, all XRISM software is written as FTOOLS and distributed through the XRISM teamat ISAS/JAXA and the NASA/GSFC HEASARC. This includes the tools used in theprocessing. All calibration files are distributed through the HEASARC caldb (calibration
database) system. This enables users to apply any calibration updates themselves. The
XRISM team at ISAS and the NASA XRISM GOF provide additional FTOOLS that may
be necessary or desirable for analyzing XRISM data. The use of other software packages
will only be supported at a lower priority level.

Chapter 4

X-Ray Mirror Assembly (XMA)XRISM has two X-ray Mirror Assemblies (XMAs) focusing X-rays onto the focal plane
detectors. One of these mirrors, Resolve-XMA, focuses X-rays on Resolve, while the other,
Xtend-XMA, is used with Xtend. Both mirrors have identical design and improve on the
mirrors on-board Suzaku. The entire system (see Figure 4.1) consists of an X-ray mirror,a Pre-Collimator (PC) stray light baffle, and a thermal shield.

Alignment Cube (AC)Thermal shield (TS)(film not shown)

PC

bla
im

Pr

Pre-collimator
(PC)

de

s

ary

rs
cto
e
l
ref

ry
da
on torsc
Se flecre203 nested
reflectors/bladesat each stage

X-ray Mirror

Interface to the S/C (top plate)Figure 4.1: Image of the XMA with all its components.

16

Proposers’ Observatory Guide

17

Figure 4.2: Picture of the X-ray Mirror Assemblies, without thermal shields (right: Resolve-XMA, left:
Xtend-XMA).
Table 4.1: XMA parameters and their characteristics.

XRISM XMA
Focal length
5.6 m
Effective aperture diameter
12 – 45 cm
Height
20 cm
Number of reflectors
203
Reflecting surface
gold
Grazing angle
0.15◦ – 0.57◦
Energy range1
0.3 – 15 keV
Effective Area2587 cm2 1.5 keV
438 cm2 4.5 keV
418 cm2 6.4 keV′
Angular Resolution (HPD)
1.3 for Resolve-XMA and 1.5′ for Xtend-XMA
Angular Resolution (FWHM)
8′′ for Resolve-XMA and 7′′ for Xtend-XMA
1 The nominal energy range of XMAs is 0.3 – 15 keV. XMAs have sensitivity tohigher energy (> 15 keV) photons but the calibration accuracy declines. 2 Includes
telescope, pre-collimator, thermal shield but not the detectors. Values are from
Boissay-Malaquin et al. (2022).

Proposers’ Observatory Guide

4.1

XMA Components

4.1.1

Mirror Part18

The XMA mirror, with an outer diameter of 45 cm, and a 5.6 m focal length, uses conicalapproximations to the Wolter Type I configuration with 203 nested shells consisting of1624 individual reflectors (Figure 4.2). The reflective surfaces of the reflectors consist of asingle-layer coating of gold. The reflectors are arranged in four quadrants for each set ofprimary and secondary housings. Seven holding bars on the top and bottom sides support
the reflectors. The reflector substrates are composed of Aluminum, and the 203 shells areorganized into three groups, with different substrates thicknesses per group. An epoxy
buffer layer helps focus grazing-incidence X-rays with energies up to 15 keV. In addition
to the reflectors, bare substrate foils are installed inside the innermost primary reflectors
to stop the incident X-rays from directly arriving at the innermost secondary which is apath producing the stray light (see Section 4.1.2). A pre-collimator, installed for straylight protection, consist of coaxially nested cylindrical aluminum blades placed above eachreflector. A thermal shield (TS) is attached in front of the pre-collimator to stabilize thethermal environment of the XMAs. The shield is made of Al-coated Polyimide to ensure alarge effective area in the soft energy band. The thermal shield transmits soft X-rays but
helps to reduce the radiative heat transfer into and out of each XMA.
The final result is a relatively light mirror with large throughput over a broad energy range (Serlemitsos et al., 2010, and references therein). Although largely inspired by
Suzaku’s mirrors, the two XMA mirrors benefited from both relaxed weight constraints andseveral fabrication improvements, including thicker substrates, a larger number of forming
mandrels, thinner epoxy layer for replication, stiffer housings, and higher-precision alignment. All these improvements yield better overall mirror performances, including a smaller
angular resolution and an improved effective area both at 1 keV and at 6 keV (Okajima etal., 2012). The XMA was developed at NASA GSFC in collaboration with JAXA/ISAS,
and Tokyo Metropolitan University in Japan. The summary of each mirror performancecompared when necessary to the relevant requirement is given in Table 4.1. Unless explicitly noted, the parameters mentioned in the table are identical for both mirrors.

4.1.2

Pre-Collimator

The grazing-incidence optics adopted for X-ray imagining use a large number (203 in thecase of XRISM ) of reflectors nested as tightly as possible to increase the effective area. As aresult, this configuration increases the possibility of reflection other than the normal double
reflection within the telescope. These contaminations of abnormal paths are referred to as“stray light.” Studies on the ASCA stray light (Mori et al., 2005) showed that the main
contribution to the stray light is from secondary reflections of photons with an incident
angle between τ and 2τ (where τ is the angle of the primary reflector measured from theProposers’ Observatory Guide

19

Figure 4.3: Major reflection paths occurring in the XMA structure. (a) Normal double reflection of the
X-ray arriving from the on-axis direction. Incident X-rays are bent by the angle 4τ in total, and converge
to the on-axis focus: (b) Secondary reflection, which arrives at the focal plane only if the incident angle ofthe X-ray, θ, measured from the optical axis is in the range τ < θ < 2τ : (c) Stray light path that gives
rise to the brightest ghost among various backside reflections, the reflection at the backside of the primary
followed by the normal double reflection. This pattern occurs when the X-ray incident angle is in the range
2τ < θ < 3τ . From Mori et al. (2005).

optical axis of the telescope – see Figure 4.3). Therefore, a pre-collimator is installed inorder to address the issue of both secondary and backside reflection that would otherwise
contaminate data on the focal plane.
The Pre-Collimator (PC) built for XRISM is derived from the one on Suzaku. EachPC consists of cylindrical aluminum shells (blades) with varying radii of 60 – 225 mm,alignment frames to guide the blade positions, and the blade housing body. Each PC bladeis placed precisely on top of the respective reflector to reduce off-axis X-ray photons that
leads to a “ghost” image within the detector field of view. Since the secondary reflection iscaused by the off-axis X-rays going just above the primary mirrors, these cylindrical blades
can block this component effectively. The blade thickness is designed to be narrower thanthat of the reflector. Thus, there is no loss of the on-axis effective area if the alignment
between the mirrors and blades is attained. The alignment frame and the housing are
made of Aluminum. Heat-forming process is introduced to the production to stabilize theblade shape in orbit. Precise curvature of radius and the linearity along with the directionof incident X-rays ensure that the blades do not obscure the telescope aperture. Figure 4.4
shows a pre-collimator and its geometry and Table 4.2 shows the main characteristics ofthe XMA pre-collimator.

Proposers’ Observatory Guide

20

Table 4.2: The main characteristics of the XMA pre-collimator.

Number of segments
Material
Number of blades
Number of support bars

4
Aluminum
204
7

Figure 4.4: Image of a Pre-collimator quadrant used for XRISM XMA.

Figure 4.5: Image of a thermal shield quadrant used for XRISM XMA.

Proposers’ Observatory Guide

4.1.3

21

Thermal shields and heaters

Per requirement, the operating temperature of the telescopes in orbit must be maintained
at 22◦ ± 7 ◦ C. This is achieved by covering the entrance side with a thermal shield (TS)composed of the 4 thermal shield quadrants shown in Figure 4.5 for each XMA. Theseshields isolate the telescopes thermally from space and reflect infrared radiation from theinterior of the spacecraft. The thermal shields also work to block optical light coming
either from the sky or from the surface of Earth illuminated by the Sun.
The shields are made of an aluminized Polyimide glued to a supporting stainless-steel
mesh. The reflectivity in the optical wave band is more than 90%, and yields a rejectionof the optical light from the bright Earth and the reflection of the IR radiation from thetelescopes at ∼290 K back to the interior of the satellite. In addition to the shields, heaters
are affixed to the wall of the housing to keep it within the mandated temperature range.
The whole XMAs are inside the sunshade base. Thermometers are attached to the wall tomonitor the temperature and control the heaters.

4.2

Expected Performance

4.2.1

Focal Positions and Optical Axis

For each mirror, the transmission is maximized when the target is observed along theoptical axis. The left diagram of Figure 4.6 shows the “ideal” telescope+detector system.
The actual configuration of a mirror includes a shift between the focal point and thedetector aimpoint (shown in the middle diagram) and a tilt of the mirror optical axis withrespect to the nominal aim-point (shown in the right diagram). Moving off the detector
aimpoint from the mirror on-axis focal point results in reducing the effective area due tovignetting. This places the boresight measurement and the determination of the opticalaxis position as one of the highest priority among the calibration requirements for both
instruments. The detector aim point is defined as the center of the Rosolve field of view.
The optical axis is tilted about 15” with respect to the nominal aim-point direction, as
this angle makes the maximum on-axis effective area. With this small angle, the loss ofthe effective area is less than 0.5% due to vignetting and is negligible. However, the focal
point may shift to +Y direction as indicated in Figure 4.7. Xtend is on-axis at aim point(less than 3” shift). Detailed in-flight calibration is forthcoming. However, the alignment
configuration measured at the ground should be adequate for Cycle 1 proposal preparation.
Figure 4.7 indicates the location of the detector aim point with respect to the Xtend andResolve FoVs. The boresight and the pointing accuracy are around 20 arcsec for Resolve
and Xtend.

Proposers’ Observatory Guide

22

Figure 4.6: Cartoon view of the actual configuration of a mirror compared to the ideal geometry (shownon the left). The middle picture illustrates the shift between the focal point and the detector aim point
and the right picture shows in addition the effect of a tilt of the optical axis with respect to the nominal
aim-point direction (HITOMI in-flight calibration plan, 2015.)

4.2.2

Effective Areas

The predicted effective area of both the Xtend-XMA and Resolve-XMA can be calculatedusing ground measurements combined with ray-tracing simulations. Table 4.3 lists theground measurements for both the Xtend-XMA and Resolve-XMA (Boissay-Malaquin etal., 2022). The simulated effective area curves are plotted in Figure 4.8.
Table 4.3: Effective area (in cm2 ) for both XMAs (with 68% statistical errors), at different energies E,
for the whole telescope axis. From Boissay-Malaquin et al. (2022)

Energy [keV]
1.5
4.5
6.4
8.0
9.4
11.0
17.5

Effective Area (cm2 )Resolve-XMA Xtend-XMA
584.7 ± 0.4
589.4 ± 0.4
434.7 ± 0.6
441.5 ± 0.6
416.0 ± 0.6
422.2 ± 0.6
345.3 ± 0.8
349.2 ± 0.8
233.4 ± 0.6
235.5 ± 0.6
163.4 ± 0.4
164.5 ± 0.4
38.4 ± 0.2
37.9 ± 0.1

Proposers’ Observatory Guide

23

Xtend FoV (a look-up view)

SAT X

Resolve FoV

SAT Y

Optical Axis

Δx ~1-2”

Δy~15”
5.4’

Aim Point

DET Y

Pointing Accuracy

DET Y

5.0’
DET X

DET X

Figure 4.7: Left: Layout of Xtend FoV (black) superimposed with Resolve FoV (cyan). This is a look-up
view. The expected Xtend aimpoint is located from ∼5.0′ (DETX) and ∼ 5.4′ (DETY) from the edges of
the CCD. The ground calibration reported that Xtend is on-axis at aim point (less than 3” shift). Right:
The expected aimpoint of Resolve as the center of Resolve FoV. The ground calibration reported that
optical axis of the XMA-Resolve is slightly shifted. In-flight calibration is on-going.

600
500
400

Area (cm2)

Area (cm2)700

Resolve
Xtend

300
200
100
0

100
2

4

6
8
Energy (keV)

10

12

70
0.3

Resolve
Xtend

1Energy (keV)

10

Figure 4.8: Effective area versus energy, for Xtend-XMA (black) and Resolve-XMA (red) plotted on
the linear scale (left) and on the log scale (right). Note that these simulated curves are based on groundmeasurements and the XMA-only effective areas, thus not including QE of the detectors, etc..

Proposers’ Observatory Guide

4.2.3

24

Vignetting

The vignetting measures the variation of the effective area as a function of off-axis angle.
It also depends on the photon energy as well as the azimuthal angle. The left panelof Figure 4.9 shows the vignetting curves measured on the ground at different energies(shown by open squares on the plot) fitted with 1D Lorentzian model, for Cross direction
1 (azimuthal angle: 0-180◦ , see the right panel of Figure 4.9), for Xtend-XMA. The similar
curves for Resolve-XMA is shown in Figure 4.10, using a smaller off-axis angle range and
1D Lorentzian model. Figure 4.11 shows these vignetting curves at 6.4 keV for 8 different
azimuthal angles for both well fitted with Resolve-XMA and Xtend-XMA, revealing a very
similar mirror FoV (defined as FWHM at 6.4 keV of 1D Lorentzian model) of about 14
arcmin for both XMAs.
Figure 4.12 shows the size of XMAs’ Field of View (not the detector FoVs), as a functionof energy.
Xtend-XMA vignetting curve (Roll angle: 0 -180°),
t 1D Lorentzian

Q3
-Y

-X

225°
270°

DET Y

180°

DET X
Q4

315°

O -axis

135°

0’

Q2

6’
0°

+X

ror
Mir ll
o
R
45°
Q1

fi

90°

Xrays

+Y

SAT Y

SAT X

ff

Figure 4.9: Left: Vignetting curves of all energies, for Xtend-XMA, for azimuthal 0-180 deg, fitted withthe sum of Lorentzian and Gaussian models to the data measured on the ground. Right: Look-down
illustration of the mirror roll angle (different from the spacecraft roll angle) and the off-axis angle usedduring the ground calibration measurements (see Boissay-Malaquin et al. (2022)). The coordinates of thesatellite and detector are shown, as well as the boundaries of the XMAs’ quadrants.

4.2.4

Angular Resolution

The Point-Spread Function (PSF) was measured at nine different energies. Figure 4.13
shows the on-axis focal plane image of the Resolve-XMA and Xtend-XMA, illustratingthat the PSF has structures. The PSF structures may have to be taken into account when

Proposers’ Observatory Guide

25

Resolve-XMA vignetting curve (Roll angle: 0 -180°),
t 1D Lorentzian

Figure 4.10: Vignetting curves of all energies, for Resolve-XMA, for roll angle (azimuthal angle) 0-180
deg, fitted with 1D Lorentzian model fitted to the data measured on the ground.
fiVignetting curve of Fe Kα, at di erent roll angles,
for Xtend-XMA

ff
ff

Roll angle 315 - 135°
Roll angle 45 - 225°

E ective area (cm2)

E ective area (cm2)Vignetting curve of Fe Kα, at di erent roll angles,
for Resolve-XMA

Roll angle 315 - 135°
Roll angle 45 - 225°

Roll angle 0 - 180°

Roll angle 0 - 180°

Roll angle 90 - 270°

Roll angle 90 - 270°

O -axis angle (arcmin)
ff

ff
ff

O -axis angle (arcmin)ff

Figure 4.11: Vignetting curves at 6.4 keV fitted to the data points measured on the ground, for ResolveXMA (left) and Xtend-XMA (right), for 8 roll (azimuthal) angles (see the right panel of Figure 4.9).

planning to observe a bright source or analyzing extended sources with a bright point source(Chapters 7 and 8). One-dimensional PSF, which was normalized so that the total numberof photons is unity within a radius of 8′ , and of the Encircled Energy Function (EEF) for
both Xtend-XMA and Resolve-XMA are given in Figure 4.14 and Figure 4.15, respectively.
Both figures also exhibit that the profiles are a function of the incoming photon energy.
Table 4.4 summarizes the values of the Half-Power Diameter (HPD) angular resolution. For

26

FoV (arcmin)

Proposers’ Observatory Guide

Energy (keV)

Figure 4.12: Field of View of the XMAs, as a function of energy. The FoV is defined here as the FWHM
value determined by fitting the vignetting curves with a 1D Lorentzian model, and does not take thephysical size of the detectors into account. The FoV is the average value of the FWHM in all directions.
The data points are from the ground measurements.

both mirrors, the measured HPD is found to be better than the 1.7′ requirement (Tamura
et al., 2022). The typical error of the angular resolution is about 0.01′ in HPD, based on
statistical error, uncertainty in the position of the center of the image, and the fluctuationof the CCD background. The angular resolution of the Resolve-XMA is about 0.1′ -0.2′
better than that of Xtend-XMA at all energies.
Table 4.4: The HPD obtained at seven energies. From Tamura et al. (2022)

Energy [keV]Resolve-XMA [′ ]Xtend-XMA [′ ]

4.2.5

1.49
1.29
1.47

4.50
1.30
1.46

6.40
1.30
1.47

8.05
1.28
1.47

9.44
1.22
1.40

11.07
1.19
1.37

17.48
1.24
1.37

Stray Light

Stray light designates X-rays that enter the detector via a path other than the normal
double reflections. It includes the case of a single reflection on either reflector, a reflection
on the back side of the reflectors, or a reflection on the pre-collimator blade. The precollimator significantly reduces the amount of stray light (as seen in 4.1.2), but it also
causes slight stray light. Stray light that can occur in a multi-nested thin foil optic was
discussed in detail by Mori et al. (2005). Figure 4.16 shows the stray light images at
1.49 keV obtained at 30′ and 60′ off-axis angles at an azimuthal angle. In both images,
the X-ray source is positioned on the right side. This is the same stray light structurereported for the Hitomi SXT mirrors (Iizuka et al., 2018). The 30′ off-axis image has
a component biased to the right side of the image, while in the 60′ off-axis image thereProposers’ Observatory Guide27

is stray light spread over the entire detector. The stray light seen on the right side of
the 30′ off-axis images is due to the reflection of the pre-collimator, while the extended
component is mainly due to the reflection on the backs of the reflectors. We note thatthe direct component coming through the gap between the inner wall of the housing andthe innermost reflector reported for Hitomi SXT Mirror (Iizuka et al., 2018) is not seen in
XMAs, because bare aluminum foil was added inside the innermost reflector to removes the
direct incident component (Tamura et al., 2022). The main contributions on the measured
stray light for XMAs are probably due to pre-collimator reflections. The amount of straylight entering within an 8′ radius, converted to effective area, is shown in Table 4.5 aswell as the amount of stray light entering the Resolve detector. We note that the flux ofthe stray depends on the azimuthal direction of the incident photons, and the values inTable 4.5 are instances. The detailed analysis of stray light will be performed with thein-flight calibration data. However, expected stray light effects will not be large basedon Suzaku observations. Therefore, proposers are not expected to consider stray light
contamination on their targets when planning proposal submissions.
Table 4.5: The effective area at 1.49 keV for off-axis sources due to stray light. From Tamura et al. (2022)

30′ off-axis [cm2 ]60′ off-axis [cm2 ]Resolve-XMA (r<8′ )0.26 ± 0.01
0.18 ± 0.01

Resolve-XMA (Reslove FoV)
0.007 ± 0.001
0.009 ± 0.001

Xtend-XMA (r<8′ )0.16 ± 0.01
0.11 ± 0.01

Proposers’ Observatory Guide

28

Resolve-XMA

DET Y

DET X

Xtend-XMA

DET Y

DET X

Figure 4.13: Focal plane images of Resolve-XMA (top) and Xtend-XMA (bottom) at the six energies of
1.49, 4.50, 6.40, 8.05, 9.44, and 11.07 keV (Tamura et al., 2022). The image size is 6.7′ × 6.7′ .

Proposers’ Observatory Guide

Normalized surface brightness

8.05 keV
9.44 keV
11.07 keV

0.1
0.01
10−3
10−4
0

2

4
Radius [arcmin]

1.49 keV
4.50 keV
6.40 keV

1
Normalized surface brightness

1.49 keV
4.50 keV
6.40 keV

129

68.05 keV
9.44 keV
11.07 keV

0.1
0.01
10−3
10−4

8

0

2

4
Radius [arcmin]

6

8

Figure 4.14: One-dimensional PSFs of Resolve-XMA (left) and Xtend-XMA (right). The six PSFs at
1.49 (red), 4.50 (green), 6.40 (blue), 8.05 (light blue), 9.44(violet), and 11.07 keV (dark blue) are plotted(Tamura et al., 2022). These plots are based on the ground measurements.

11
1.49 keV
4.50 keV
6.40 keV
8.05 keV
9.44 keV
11.07 keV

0.6
0.4
0.2
0

0

5
10
Diameter [arcmin]

1.49 keV
4.50 keV
6.40 keV
8.05 keV
9.44 keV
11.07 keV

0.8
Normalized Intensity

Normalized Intensity

0.8

15

0.6
0.4
0.2
0

0

5
10
Diameter [arcmin]

15

Figure 4.15: The EEF of the Resolve-XMA (left) and the Xtend-XMA (right). The six EEFs at 1.49
(red), 4.50 (green), 6.40 (blue), 8.05 (light blue), 9.44 (violet), and 11.07 (dark blue) are plotted (Tamura
et al., 2022). These plots are based on the ground measurements.

Proposers’ Observatory Guide

30

Figure 4.16: Stray light images of Resolve-XMA (top) and Xtend-XMA (bottom) at 30′ off axis angle(left) and 60′ off axis (right) (Tamura et al., 2022). The X-ray energy is 1.49 keV. The image size is 17.8′
×17.8′ . The FOV of Resolve is shown in left top panel.

Chapter 5

Resolve
Resolve is a high resolution X-ray imaging spectrometer with a 3′ × 3′ field-of-view. It
consists of X-ray focusing mirrors (XMA, Chapter 4), and a 36-pixel X-ray calorimetersystem (XCS) with better than 5 eV resolution over much of the 1.7-12 keV energy band,
according to the initial in-orbit calibration. An anti-coincidence detector placed behindthe calorimeter array enables rejection of particle background events.
Key performance parameters are provided in Table 5.1 and an overall view of the
Resolve is shown in Figure 5.1. Figure 5.2 shows the effective area of the instrument.
The nominal energy range of Resolve is 1.7-12 keV. At low energies, the effective areaof Resolve declines steeply due to the closed gate valve (∼0.5 cm2 at 1.7 keV). At highenergies, Resolve is sensitive to photons above 12 keV, even up to ∼20 keV. However,calibration becomes less reliable at high energies, and 12 keV is currently the high end of
the nominal Resolve energy range.
Resolve is descended from the X-ray Spectrometer onboard Suzaku Observatory (Kelley
et al., 2007) and the Soft X-ray Spectrometer (SXS) onboard Hitomi (Kelley et al., 2016).
It is the first microcalorimeter detector that is available to guest observers.

5.1

Overview

Figure 2.2 shows the schematic of the position of the Resolve within the satellite. X-rays
emerging from the X-ray Mirror Assembly (XMA; see Chapter 4) pass through a Filter
Wheel Mechanism (FWM; cf. Section 5.8.1) mounted 91.5 cm from the detector. The
FWM has six filter positions: two open, one with 55 Fe calibration source, two for X-ray
attenuation to reduce the flux incident on the detector, and one with a polyimide filter.
Not all filters are available to observers (cf. Section 5.8.1). Modulated X-ray Sources(MXS; cf. Section 5.6), mounted on the filter wheel axis beyond the FWM with respect
to the path of X-rays, serve to calibrate the detector gain. The filter wheel position is31

Proposers’ Observatory Guide

32

Table 5.1: Summary of the expected Resolve detector properties based on the ground calibration. HPD
stands for Half Power Diameter.

Parameter
Lifetime
Energy range
Effective area
Angular resolution
Energy resolution
Energy-scale accuracy
Line-spread function accuracy
Pixel size
Field of View
Array format

Relative timing accuracy
Operating temperature
Maximum X-ray count rate
Residual background

Expected
Goal: 5 years
1.7-12 keV∼ 180 cm2 @6keV
1.3′ HPD
5 eV (FWHM)
0.5 eV
Goal: 1 eV
30”×30”818 µm× 818 µm
3.1′ × 3.1′
6×6
35 science pixels
1 calibration pixel
80 µs for high- and mid-res events
50 mK
200 counts s−1 array−1
0.8 × 10−3 counts s−1 keV−1

user-selectable, while MXS is not.
Beyond the FWM, the X-rays go through the aperture door (i.e. the closed gate valve)
and then enter the dewar (a cryogenic vacuum vessel where the microcalorimeter array andanti-coincidence detector reside) through five thin-film filters, each staged at a different
temperature, before reaching the detectors. These filters help isolate the cooling systemfrom the ambient environment and attenuate the flux of IR photons (thermal radiationfrom warmer structures) and UV/optical photons (from the sky) that reach the detector.
The optical blocking filters in the dewar are fixed in place. The transmission curve for thefull five-filter stack is shown in Figure 5.3.
The microcalorimeter detector array, along with the anti-coincidence detector, is housed
within the Detector Assembly and maintained at its operating temperature of 50 mK by
a multi-stage cooling system with a temperature stability better than 2.5 µK, and withefficiency >95% for more than three years. This low operating temperature and its stability
enables the detector to attain its required energy resolution. The cooling system is designedwith redundant features to be tolerant to some failure cases and to enable operation of thedetector beyond the nominal three-year lifetime.
Analog signals from the Detector Assembly are amplified and digitized. The digitizedsignal is sent to the Pulse Shape Processor (PSP), which detects X-ray events from thedata stream and applies an optimal filtering algorithm to determine the pulse heights ofProposers’ Observatory Guide

33m
00 m 4'/mm
f = 56scale: 0.61

North

DETY ft -Y)
cecra Rollap(S
4 .9 7 4

3.05')
mm (

23
21

East

20

19

10

9

13

11

14

15
16

size
rber
Abso 0.004 mm±
41
0.8 (0.50')

Decl.

1

2

0

17

DETX ft +X)
cecra
(Spa

27

28

35

18

29

31

33

25

22

30

32

34

26

24

Plate

4

7
8
G ap
mm
0.018 1')
(0.0

West
3
5

6

7 .0
34
(4.3 mm2' )Pitch mm
0.832 1')
(0.5

South

R.A.

Figure 5.1: A schematic of the Resolve microcalorimeter array focal plane. Dimensions shown are the
Look-up map of the SXS detector
design values, but most actual absorbers were 0.818 mm wide, and the gaps were correspondingly smaller(Pixel 12 is the cal pix and is not shown in this figure.)
on average.
200

Effective Area [cm2]

175
150
125
100
75
50
25
0

2

4

6
8
Energy [keV]

10

12

Figure 5.2: Effective area of Resolve, which includes the gate valve transmission (Midooka et al., 2021).

X-ray events. Details of PSP are provided in Section 5.3.1.
The microcalorimeter detector does not trigger on optical events, but is sensitive to
optical photons. The absorbed energy contributes to noise, thus degrading the X-ray
energy resolution. However, photons outside of the passband are highly attenuated by theProposers’ Observatory Guide

34

1.0

Transmission

0.8
0.6
0.4
0.2
0.0

2

4

6
8
Energy [keV]

10

12

Figure 5.3: Transmission curve for the non-selectable five-filter blocking filter stack.

filter stack, and further by the currently closed aperture door (gate valve), and therefore
this effect should not be important.

5.2

The microcalorimeter

The Resolve detects and measures X-rays with a 6 × 6 pixels X-ray microcalorimeter array(see Figure 5.1). There are 35 active pixels in the main array and 36 readout channels.
One corner pixel is not wired and, instead, a calibration pixel out of the field of view usedwith a continuous, collimated 55 Fe source to provide one source of gain tracking for themain array. This pixel is active at all times and is designated as pixel number 12 (in lieu
of the inactive corner pixel; cf. Figure 5.1 and Section 5.6).
The pixels are arrayed on a 0.83 mm pitch with a filling factor greater than 96%. Eachpixel projects to 0.5′ × 0.5′ on the sky and the entire X-ray sensitive portion of the array
subtends a solid angle of 3.05′ × 3.05′ . It is very important to note that these pixels are
smaller than the PSF (0.5′ pixel vs 1.3′ HPD). This means that if there are bright sourcesin and around the FOV or the source is extended, photons from other regions of the sky
will scatter into an individual pixel. For more information, see Chapters 7 and 8.
A microcalorimeter detects X-rays using a high-sensitivity thermometer that measures
the heat produced by thermalization of photon energy when a photon is absorbed by a lowheat capacity absorber.
Since the amplitude of the signal translates into X-ray energy, a microcalorimeter is asingle photon detector that measures the energy of every incident photon. By combining
36 independent microcalorimeters in a 6×6 array configuration, the Resolve instrument
provides high-energy resolution imaging spectroscopy of point-like and diffuse celestial Xray sources without the use of dispersive optics Mitsuda et al. (2010, 2014), in contrast to
the grating instruments on board of XMM-Newton and Chandra. Unprecedented energy

Proposers’ Observatory Guide

35

resolution is attained at E > 2 keV (including the Fe K band) for all sources, and at allenergies for extended, diffuse sources such as supernova remnants and galaxy clusters.

5.2.1

Photon detection

The concept underlying an X-ray quantum microcalorimeter detector is based on the measurement of the temperature increase of an isolated thermal mass. The measurement ismade using a thermometer in a quasi-static equilibrium with the absorber (Figure 5.4).
Thus, the three essential components are (i) an absorber that absorbs the incident X-ray
and thermalizes the deposited energy, (ii) a coupled thermometer that measures the temperature increase in the absorber, and (iii) a weak link to a heat sink that restores theabsorber to its original temperature. The energy resolution is determined by the precisionof the temperature increase measurement against a background of temperature fluctuations.
For an absorber heat capacity Cpix , thermal link conductance Glink , and heat sinktemperature Tsink the absorber temperature increases by ∆Tpix ≈ EX−ray /Cpix (on theorder of a ∼mK for Resolve) upon absorption of energy EX−ray and re-equilibrates to its
quiescent temperature with a decay time constant tdecay ≈ Cpix /Glink . This is illustrated
in Figure 5.4.
For Resolve, the rise time for the temperature pulse induced by absorption of an X-ray
is on the order of a ms, and tdecay is a few ms. However, the detector requires multiple
decay timescales to fully re-equilibrate so that for sufficiently bright sources a second pulse
may arrive on the tail of a previous pulse. This is explained further in Section 5.3.3.

5.2.2

Anti-coincidence detector and particle background

Resolve is equipped with an anti-coincidence (anti-co) detector. Its main purpose is toreject cosmic-ray events that could be incorrectly interpreted as X-rays. Unlike X-ray
photons (which are detected by the microcalorimeter only), cosmic rays trigger events inboth the microcalorimeter and the anti-co, and thus they can be identified and removed.
Additionally, the anti-co provides an independent monitor of the particle environment. In
addition to the advantage of the spacecraft being placed in a low-Earth orbit, a lower andmore stable particle background compared to other X-ray missions is achieved thanks tothe anti-co.
The Resolve anti-co detector is essentially identical to that deployed in SXS on Hitomi(Kilbourne et al., 2018b). It is located directly below the microcalorimeter array, as shownin Figure 5.5. As ionizing charged particles pass through the anti-co, unbound electron/holepairs are liberated and their drift under the influence of an electric field creates a current
that is amplified and measured. Calorimeter events that arise from ionizing particles are
flagged for subsequent vetoing as they also trigger a pulse in the anti-co detector. This

Proposers’ Observatory Guide

36

Thermistor
X-ray photon

Absorber(Heat capacity, C)

Thermal link

Heat sink(< 0.1 K)

Figure 5.4: The concept of an X-ray calorimeter is shown in the left panel, whereby an X-ray deposits
energy in an absorber that is measured by a thermometer after the energy thermalizes. Subsequently, theabsorber returns to its quiescent temperature and the system is ready to detect and measure the next X-ray.
Right panel shows a schematic of the temperature of an absorber, which is a part of Resolve pixels, as afunction of time. Upon absorbing the photon, the temperature goes up proportionally to the energy of theincident photon, and then goes down with time constant proportional to the heat capacity and thermal
link conductivity.

Figure 5.5: Schematic arrangement of the XRISM Resolve microcalorimeter array and the anti-coincidence
detector (c.f. Porter et al., 2010). The microcalorimeter array sits on top of the anti-co detector.

way the anti-co detector reduces the non X-ray background.
The anti-co incidence rejection (i.e., removal of cosmic-ray events) is performed on the

ADC units

37

ADC units

Proposers’ Observatory Guide

0.00

0.02
0.04
Time [s]

0.06

0.00

0.02
0.04
Time [s]0.06

Figure 5.6: Typical shape of a pulse triggered by an incident X-ray. Left panel shows a single photon,
while right panel shows a situation where second photon is detected before the pixel equilibrates after the
first event. The undershoot following the pulse happens because the electronics are AC coupled. For details,
see Section 5.3.1.

ground as part of pipeline processing, allowing for more flexibility in applying different
veto timing windows, as well as the capability to conduct spectroscopic analysis of anti-co
events.

5.3

Event processing system

5.3.1

Event triggering

Signals from both the microcalorimeter and the anti-co detector are transferred from the
Detector Assembly to the XBOX (X-ray amplifier BOX), where they are shaped (via analog
high-pass and low-pass filters), amplified, and digitized. The digitized signals are sent to
the Pulse Shape Processor (PSP) for event discrimination, classification, and processing.
The PSP triggers on pulse candidates and characterizes them, calculating the pulse heights,
arrival times, and other characteristics of each event.
A temperature increase in the microcalorimeter due to absorption of an X-ray leads
to a drop in resistance and corresponding drop in voltage across the thermistor. The
processing results in a profile where peaks are followed by deep, long undershoots that area result of AC-coupling in the electronics. When the calculated derivative of this amplitude
profile exceeds the threshold, an X-ray event is triggered. The characteristic profile of theamplitude of this drop after shaping and inversion in the XBOX is shown in Figure 5.6.

Proposers’ Observatory Guide

38

The PSP calculates a running boxcar derivative of the incoming data stream of eachcalorimeter channel, and when that derivative exceeds the selected threshold, an X-ray
event is provisionally triggered. Then, a search for events occurring after the first event isconducted. To detect second order pulses (pulses arriving during the decay of the previous
pulse), the same trigger algorithm is applied to an adjusted derivative profile (Boyce et al.,1999). The adjusted profile is triggered as a second order pulse if the adjusted derivativerises above a threshold and then falls below zero within a specified length of time. The
latter condition assures that the second order pulse has the characteristic shape of an X-ray
event (Figure 5.6, right panel). The process is repeated to search for third order events,and so on until no further pulses are found.
Anti-coincidence (anti-co) events (cf. Section 5.2.2) are triggered in a similar way to
pixel pulse events, but are based on the pulse profile only (and not its derivative). Anti-co
events are utilized to flag coincident pixel events for subsequent filtering. Triggering occurs
when the sample value exceeds the inter-pulse baseline level by a sufficient amount.

5.3.2

Pulse height determination

To determine the pulse height amplitude (PHA; which, ultimately, is related to the energyof the incident photon), the signal processing by the PSP exploits the approximate energy
independence of the pulse profile and energy resolution by employing the optimal filtering
method. This requires the construction of a pulse profile template. Cross-correlation of
the observed signal with this optimal filter template yields an estimate of the pulse heightwith the lowest signal-to-noise frequencies filtered out. The details of this method can befound in Moseley et al. (1988) and Szymkowiak et al. (1993).
Since the needed precision of the energy determination of anti-co pulses is much lessthan for calorimeter pulses, to conserve the computing resources of the PSP, the assigned
anti-co pulse height is simply the peak value of the pulse.

5.3.3

Event grading: low, mid, and high-res grades

To assure the highest accuracy of the optimal filter matching (as described in Section 5.3.2)
there must be only one pulse per pulse time interval. If this is the case, the highest energy
resolution (“hi-res”) is attained. If there is a second pulse within that interval, a truncatedtemplate is applied, and a lower resolution (“mid-res”) is set. If there are multiple pulses
within the truncated template the resolution is even lower (“lo-res”), and the pulse heightmust be estimated from the height of the peak over the baseline. Figure 5.7 shows thedifference in these spectral resolutions.
Specifically, the events are graded based on comparing the time difference between
neighboring pulses against two limiting values, δt1 = 70.7 ms and δt2 = 18.3 ms. Hi-res
events are those for which pulse height analysis using the full template is possible, i.e. thoseUnnormalized counts/bin

Proposers’ Observatory Guide

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39
high-resmid-res
low-res

5880
5900
Energy [eV]5920

Figure 5.7: Comparison of the fluorescent Mn Ka spectra for different event grades from the Resolve
ground calibration program.

with no other detected pulses with arrival times nearer than δt1 . Mid-res events are thosefor which pulse height analysis using the truncated template is possible, i.e. those whereother pulses are detected with arrival times nearer than δt1 but not within δt2 . If pulses
are detected too close (< δt2 ) for the mid-res template, they are classified as low-res.
Events that are separated in time from preceding events by a full template length ormore are characterized as primary events, otherwise as secondary events (primary events
may then be classified as either high-resolution, mid-resolution, or low-resolution based on
the time interval to the subsequent event). Secondary events are those that are precededby an event within the time interval allocated for a single pulse sample, and classified as
mid-res or low-res depending on the time interval to the nearest event using these same
criteria.
Because the detector may not be fully re-equilibrated from the energy absorbed fromthe previous pulse there may be a shift in the gain for secondary events (cf. Section 5.4).
For high-count rate sources, there may be a long chain of secondary events associatedwith a single primary. Finally, if pulses arrive so rapidly that they cannot be individuallydetected at all, they are measured as a single event with an amplitude equal to the sumof those of the unresolved components. This pile-up will result in a biased spectrum andcount rate. Note that this happens only at very high count rates (cf. Chapter 8).
Combining high, mid, and low resolution grades (based on the relative time of thenearest event) with primary and secondary flags leads to the five possible event gradeclassifications illustrated in Figure 5.8 and shown in Table 5.2. The Resolve spectral
resolution requirement applies to Hp and Mp grades (marked in bold font in Table 5.2),
though we expect Ms to be very similar.
All detector pixel pulse events are therefore graded as a combination of primary orProposers’ Observatory Guide40
δt1

Primary event
Secondary event

δt2
Hp eventMp event

Ms event

Pulse heightLp eventLs eventTime

t0

Figure 5.8: Sketch illustrating the grading of Resolve events. The five cases described in the text (highprimary, medium primary, medium secondary, low primary, and low secondary) are considered. For details
see Ishisaki et al. (2018).
Table 5.2: Definition of event grades. Calorimeter-grade energy resolution is obtainable only for events
labeled Hp and Mp, marked in bold fonttn ≤ δt2
δt2 < tn ≤ δt1
δt1 < tntp ≤ δt2LSLSLS

δt2 < tp < leδt1
LS
MS
MS

δt1 < tp
LP
MP
HP

secondary and high, medium, or low resolution – depending on the proximity to adjacent
events in time. The event grade provides a measure of data quality.

5.3.4

Event grade branching ratios

The proportion of events in the various event grades depends on the source count rate.
The highest-resolution data is only obtainable at count rates that are sufficiently low suchthat consecutive events are well separated in time.
The branching ratio, defined as the fraction of events that fall into a particular event
grade, may be calculated from the template width assuming Poisson statistics. The ratio
as a function of the event rate landing on a single pixel is shown in Figure 5.9 for each
grade and discussed further in Chapter 8.

Proposers’ Observatory Guide

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Figure 5.9: The event grade branching ratio is changed by the incoming rate. At low rates, most events
will achieve calorimeter resolution (Hp or Mp ). For bright sources, resolution is degraded.

As the rate increases, the branching ratio for calorimeter-resolution (i.e., HP +MP
grade) events declines. The HP +MP event rate reaches its maximum of 3.6 events s−1
at an incoming rate of about ∼10 counts s−1 pixel−1 .

5.3.5

Pulse Shape Processing limit

There is a limit to the number of events that can be processed by the PSP in a given time:
count rate of ∼200 s−1 array−1 . When the count rate is high and this limit is exceeded,some pulses will not be processed and dead time arises. The dead time can be calculatedby the PSP telemetry and should be carefully taken into account.
The computational load is the limiting factor for the PSP to process high-count rate
sources, the limiting rate depending on the event grades. Choosing an appropriate filter
may help increase efficiency of calorimeter-grade events (see Section 5.8.1).

5.3.6

Cross-talk

Electrical and thermal cross-talk can cause further degradation of the energy resolution.
Electrical cross-talk in practice is only a concern between nearest-neighbor electrical channels. Most electrical cross-talk from in-band X-rays is not triggered, but these un-triggered
pulses can contaminate X-ray pulses. Additionally, the heat flow into the pixel frame froma photon absorbed by a calorimeter pixel or a photon directly absorbed in the frame causes
a temperature perturbation, which also contributes to noise at very high count rates. For
more details, see Chapter 8.
This cross-talk can significantly impact the energy resolution for high count rate ob-Proposers’ Observatory Guide

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servations. Ground tests demonstrated energy resolution degradation of a few eV. Theenergy resolution can be recovered by screening out events that occur simultaneously onelectrical nearest neighbors, in which case the cross-talk cut is applied at the price of lower
effective exposure time (note that this approach is not applicable when e.g. dead timearises due to a high count rate, cf. Sect. 5.3.5). This was a consideration for the Resolve
instrument team in planning the usage of the MXS but it will not be a significant effectfor most observations.

5.4

Detector gain

The PSP measures the pulse height of a detected event. This pulse height is then converted
into incident photon energy using the detector gain. The latter is determined by theeffective temperature of each pixel, which depends on both the heat-sink temperature andradiative loading. The effective temperature is determined by using a reference X-ray line
(nominally from the FW 55 Fe source and/or MXS) and performing non-linear interpolation
between calibration curves taken at different heat-sink temperatures. The gains can becontinuously calibrated or interpolated between intermittent assessments, depending on
the time scale of the change.

5.5Line spread function

The Line Spread Function (LSF) describes the response of the detector to a monochromatic
beam of photons. A pre-flight LSF has been characterized for each pixel, as shown in
Figure 5.10. It consists of a Gaussian core, which dominates up to a level of ∼1% for highenergies. A minor fraction of events will be redistributed to lower energies in an exponential
tail, electron-loss continuum, and escape peaks (Eckart et al., 2018). The LSF depends onthe detector noise, and therefore these pre-flight measurements will be repeated in-orbit.

5.6

Calibration sources and MXS operations

Good calibration is critical for a high spectral resolution detector like Resolve. Both LSF(Section 5.5) and gain scale (Section 5.4) need to be regularly calibrated for each pixel.
Resolve is equipped with three types of calibration sources to monitor the temporal changeof the detector gain in the orbit (gain drift). One is the calibration pixel (cal-pixel, or pixel
number 12) that does not receive celestial X-rays but instead is constantly illuminated bya collimated 55 Fe radioactive source. Another 55 Fe source is mounted on one of the filter
positions of the Filter Wheel (cf. Section 5.8.1) and can be used to illuminate the wholearray when the position is selected. Mn Kα events at the cal-pixel are always available forProposers’ Observatory Guide

43

105
Counts [s−1 keV−1]

Counts [s−1 keV−1]

105
104
103
102
101

104
103
102

5.38

5.39

5.40
5.41
Energy [keV]

5.42

5.43

101
5.20

5.25

5.30

5.35
5.40
Energy [keV]

5.45

5.50

Figure 5.10: Instrument response to a monochromatic emission line at 5.414 keV, demonstrating theGaussian core which dominates the line shape and the electron loss continuum. For more details, see
Eckart et al. (2018)calibration reference for the overall gain drift of the array (common-mode gain), while thefilter wheel 55 Fe is useful to monitor the gain drift pixel by pixel (differential-mode gain).
The last type of calibration sources are the Modulated X-ray Sources (MXS), which
can be used to monitor both the common and differential gain drift and is not available toobservers. Four MXS units are mounted below the Filter Wheel. Two of these are nominal
and the other two are redundant. Each pair consists of direct and indirect (fluorescent)sources that emit a set of high-energy X-ray emission lines, mainly Cr Kα at 5.41 keV andCu Kα at 8.05 keV (de Vries et al., 2018).

5.7

Background

There are multiple contributions to the Resolve background. Astrophysical backgrounds
and foregrounds include e.g. the Cosmic X-ray Background (CXB), the Galactic Halo(GH), the Local Hot Bubble (LHB), and Solar Wind Charge Exchange (SWCX). The
significance of these components will depend on the position on the sky, on the timing of
the observation, and on the target itself (diffuse or not, etc). With the aperture door (gatevalve) closed, CXB is the main source of the celestial X-ray background.
The Non-X-ray Background (NXB) originates mainly from secondary particles (mainly
electrons and photons) produced by primary cosmic rays that do not traverse the anti-coand X-rays from the calibration sources (scattered from the internal 55 Fe source, or the
MXS pulse tail). Figure 5.11 shows the NXB measured in-orbit by the Hitomi SXS microcalorimeter. The Mn K lines are scattered X-rays from the internal calibration source,
and the other lines are fluorescence produced by cosmic rays. These measurements demonstrated that the NXB in microcalorimeter detectors with an anti-coincidence veto is usually
very low (Kilbourne et al., 2018a). However, we caution that the NXB can have a non-Proposers’ Observatory Guide

44

Figure 5.11: Spectrum of the NXB background measured in-orbit in the Hitomi SXS microcalorimeter.
For details, see Kilbourne et al. (2018a).

negligible impact in a few specific cases (e.g. sources with very faint, diffuse emission).

5.8

Planning observations with Resolve

It is strongly recommended that proposers simulate Resolve observations as part of theirscientific justification and feasibility. Special caution must be taken when proposing forextended and bright sources (cf. Chapters 7 and 8). An up-to-date list of simulation tools
can be found on the NASA XRISM website1 .

5.8.1

Filter selection

The filter wheel provides a means of selecting a filter setting so as to optimize observations.
The filters essentially allow the reduction of flux from bright celestial sources to reduce thedetector dead time. The six filter wheel positions include two empty positions (i.e., nofilter), a position with radioactive 55 Fe calibration source, Be and neutral density (ND)
filters for X-ray attenuation, and a polyimide filter. Filter wheel is shown in Figure 5.12
and transmission curves are shown in Figure 5.13.
There are a total of six positions, three of which can be selected by the observers:
1

https://heasarc.gsfc.nasa.gov/docs/xrism/proposals/

Proposers’ Observatory Guide

45

#5
#6
#4

#1
#3
#2
Figure 5.12: Filter wheel filter configuration arrangement, viewed looking up from the detector toward
the sky.

• Positions 1 and 4 are open, filter-less positions. Most nominal observations will be
conducted with one of the two open positions.
• Position 3 has an X-ray Neutral Density (ND) filter consisting of a 0.25 mm molybdenum foil, which blocks 75% of the flux uniformly across the 1.7-15 keV energy
band. Molybdenum does not produce fluorescence lines in the Resolve energy band.
This filter can be used to reduce the flux of bright sources, therefore improving theefficiency of microcalorimeter-grade events and, ultimately, the spectral resolution(see Sections 5.3.3 and 5.3.5 as well as Chapter 8 for details).
• Position 5 has a 30 µm Be filter used to preferentially suppress low-energy photons(below 3 keV). This filter maintains both flux and high spectral resolution at the Fe
K region.
Positions that are not selectable are:

Proposers’ Observatory Guide

46

1.0

Filter transmission

0.8
0.6

ND filter
Be filter

0.4
0.2
0.0

2

4

6
8
Energy [keV]

10

12

Figure 5.13: Transmission curves for the selectable filters, the Neutral Density (ND) and Be-filter. For
filter effective area curves see Figure 8.3.

• Position 2 has a 200 nm thick polyimide filter, fixed on a stainless steel mesh. Thisfilter is to be used for protection against contamination during the initial phase of
the mission only.
• Position 6 has the radioactive 55 Fe sources illuminating the entire array and producing emission lines at 5.89 keV and 6.49 keV. These five sources, all fixed on two
perpendicular bars across the FW opening, may be used for in-flight gain and energy
resolution calibration, in a manner that is determined by the Resolve instrument
team.

5.8.2

Observing faint to moderately bright point sources

To plan an observation of a faint to moderately bright point sources, simple spectral simulation and background modelling is sufficient to determine the integration time necessary
to accumulate sufficient number of counts.

5.8.3

Observing extended sources

We refer the reader to Chapter 7.

5.8.4

Observing bright sources

We refer the reader to Chapter 8.

Chapter 6

Xtend/SXI
6.1

Xtend Basics

Xtend is the soft X-ray imaging telescope of XRISM , with a pair of an X-ray mirror (XMA)
and an X-ray CCD imaging spectrometer (Soft X-ray imager: SXI, Hayashida et al., 2018).
Xtend has a larger field of view (38′ ×38′ ) and a low/stable non-X-ray background, ideal forobserving diffuse X-ray emission with low surface brightness. The large field of view (FoV)
and the finer pixel size (1.′′ 74) also provide more detailed spatial information of or aroundthe target, extending the capabilities of Resolve’s spectroscopy by evaluating contaminationfrom sources outside Resolve’s FoV and measuring the local diffuse background (Figure 6.1).
Xtend works in a complementary role to Resolve.
Besides a few improvements, Xtend is almost identical to the Astro-H SXT + SXItelescope system (see Section 6.4.6). Therefore, its publications and archival data should
be a good reference for Xtend’s performance. This chapter describes the SXI module andthe combined performance of the XMA and SXI. Chapter 4 describes the XMA-specificperformance, which has an identical design to Resolve’s XMA telescope but with a slightly
different performance.
The SXI houses four identical X-ray CCD sensors in a 2×2 grid. Each CCD has an
identifier, CCD ID=0−31 , used for the science and housekeeping data (Figure 6.2). These
CCDs are similar to those on earlier or ongoing X-ray observatory missions, such as Suzaku
XIS, Chandra ACIS, and XMM-Newton EPIC-MOS. However, the SXI CCDs utilize thep-channel, backside-illumination technology, which realizes a thick depletion layer of 200µm, with high quantum efficiencies in a broad energy range between 0.4–13 keV. The
following list describes the key instrument performance. Table 6.1 also summarizes Xtend’s
1

The instrument team and manufacturer also uses another identifier CCD1-4, which has the same
numbering order but starts from 1 (e.g., CCD1 for CCD ID=0). This document consistently refers to the
CCD ID identifier to avoid confusion.

47

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48

Figure 6.1: X-ray images of the Perseus cluster obtained with the Hitomi observatory (Nakajima et al.,2018). (a) Full X-ray CCD (SXT+SXI) image overlaid with the micro-calorimeter (SXT+SXS) FOV incyan. (b) Magnified CCD image of the central region overlaid with the SXS pixel boundaries in cyan. (c)Micro calorimeter (SXT+SXS) image of the same region.

characteristics.
Table 6.1: Xtend Characteristics

Field of View
38′ ×38′ (Full window)
Sensitive band
0.4 − 13 keV
Effective area
356 cm2 @1.5 keV, 307 cm2 @6 keV†
On-axis XMA PSF at 6.4 keV
1.47′ (HPD), 7.2′′ (FWHM)
Pixel size
1.74′′ (48µm, logical)
Time resolution
4 sec (Full window), 0.5 sec (1/8 window)
Energy resolution (FWHM)∼180 eV @ 6 keV
Pileup tolerance
2.5 mCrab (Full window)
Total (NXB + Sky) X-ray background‡ ≤10−6 counts keV−1 s−1 arcmin−2 cm−2
HPD: Half Power Diameter.
† See the XMA chapter (4) for the details.
‡ See Figure 6.4 right and Figure 18 of Nakajima et al. (2018) for the details.
These values are based on ground or in-flight calibrations.

Proposers’ Observatory Guide~44-59’’(25-39 pix)

18.9’ (640 pix)
CCD_ID=0

D

CCD_ID=2

18.9’ (640 pix)

(CCD1)~44-59’’(25-39 pix)

49

A

(CCD3)

Blockage
area
C

B

B

C

’o

15

A

-ax

is

’o

10

is

-ax

D

5’

o

D

is

-ax

A

5.4’
5.0’

C

Resolve FoV

B(3.05’ x 3.05’)

B

C

SAT X

DETY
Roll

DETY
A

DETX

CCD_ID=1(CCD2)

CCD_ID=3(CCD4)

55Fe calibration sources

SAT Y

D

DETX

Figure 6.2: Left — Schematic layout of the SXI CCDs looking toward the sky from the back of the focal
plane. The 4 CCD chips are aligned in a 2×2 format. Each CCD has two electrically separated segments,
whose boundary is shown by the dashed line. Each segment has two readout nodes for redundancy (theblack dots with the A/B/C/D labels). Each chip has an additional frame store area on the readout nodeside. The filled gray half circles show approximate areas illuminated by the calibration sources. The redbox shows the SXS FOV and the dotted orange circles show the off-axis angles. The distances from the
CCD edges are inflight measurement values. There is a tiny area on the far side corner from the aim point
shown in yellow, where a camera body structure blocks the view of the sky. Right — SXI CCD sensor
photo (Nakajima et al., 2020, Fig. 4). Gold-coated metal frames cover the frame store areas.

Field of View: The 2×2 CCD array covers a 38′ ×38′ square FOV with small mechanical gaps between the CCD chips. The gaps between the active pixel regions are 44′′ −59′′(1.2−1.6 mm). The aim point co-aligns with the Resolve aim point and offsets by 5 or 5.4′
from the CCD edges of CCD ID=1 to avoid the chip gap. In the 1/8 window mode, the2 CCDs (CCD ID=0, 1) have a narrow rectangular FOV with the aim point at the width
center. The long side does not touch the CCD edges, having a relatively big imaging gap
with the other 2 CCDs’ FOV (Figure 6.5, Section 6.2).
Sensitive Energy Range & Quantum Efficiency: The SXI CCD sensors have ap-type channel and n-type substrate, enabling a thick depletion layer of ∼200 µm andrealizing a large X-ray stopping power up to ∼13 keV. The CCDs are backside illumination CCDs with high transparency to soft X-rays. The SXI system also has the Optical
Blocking Layer (OBL) on the CCD surface to block optical/IR lights and the Contamination Blocking Filter (CBF) at the hood top to avoid molecular contamination build-up
and further block optical/IR lights. These components limit the soft X-ray sensitivity to∼0.4 keV. Figure 6.3 left shows the SXI CCD quantum efficiency.
Effective Area: This value primarily depends on the XMA’s effective area (see Sec-

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Figure 6.3: Left — SXI CCD Quantum efficiency measured with the ground experiments (black). The
plot also separately shows the quantum efficiency of a CCD chip combined with the transmission of the
Optical Blocking Layer (QE + OBL, red) and the transmission of the Contamination Blocking Filter (CBF,
green). Right — Effective area of the Xtend telescope, including the XMA component.

tion 4.2.2) and the CCD’s quantum efficiency. Beside, the SXI runs with the charge
injection operation through the mission, which produces 1−3 insensitive CCD rows every
80 rows and reduces the effective area by 1−3%. Figure 6.3 right shows the effective areaof the Xtend system for a point source at on-axis. The effective area can decline withradiation damage or contamination condensation in the line of sight. The instrument andscience operations teams will calibrate it regularly and update the result in the calibration
database.
Angular Resolution: Each CCD has a square imaging area with 1280×1280 physicalpixels of 24 µm length each, which are binned by 2 pixels onboard during the science
operation. The output science data have 640 “logical” pixels of 48 µm, equivalent to1.74′′ in the sky. This pixel size is significantly smaller than the on-axis PSF’s half-power
diameter (1.47′ ) and FWHM (7.2′′ ) (see Chapter 4).
Time resolution: The nominal operation (the Full window mode without the burst
option, see Section 6.2) takes a frame exposure for 3.96 seconds and then moves the obtained
data to the frame store area in 36.864 milli-second for readout. The imaging area takes
another exposure right after this transfer, so, this operation continuously obtains data every
4 seconds. The 1/8 window mode without the burst option runs this cycle every 0.5 sec.
The 1/8 window mode with the burst option uses only a fraction of each 0.5 sec exposure.
The data-collecting cycle does not increase, but users can use its narrow exposure interval
information for pulsar analysis, for example.
Energy resolution: The CCD energy resolution depends on the statistical fluctuationof charges produced by X-ray photons, the number of charges lost during the CCD chargetransfer (charge transfer inefficiency: CTI), and the readout electric noise. Normally,

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10−7

Hitomi SXISuzaku XIS−BI
Al K
Si K
Au Mα

Au Lα
Ni Kα

normalized counts s−1 keV−1 arcmin−2 cm−2

normalized counts s−1 keV−1 phys.pixel−1

10−3

Au Lβ

Au Lγ
10−8
5

Energy (keV)

10

15

Chandra ACIS−S3
ChandraACIS−I0123
SuzakuXIS−BI

10−4

XMM
MOS1

10−5

10−6

XMM PN
Hitomi SXISuzaku XIS−FIASCA SIS0

10−7

CXB

10−8

1

2

5
Energy (keV)

10

20

Figure 6.4: Left — Hitomi SXI NXB spectrum normalized by the physical pixel area. The plot also showsa Suzaku XIS-BI NXB spectrum for the same physical pixel area. Right — Hitomi SXI total (NXB + sky)background spectrum normalized by the effective area and solid angle of the sky. The plot also shows the
CXB spectrum and background spectra of other X-ray CCD instruments (Nakajima et al., 2018).

the first component, charge statistical fluctuation, is the most significant factor in the
CCD energy resolution. Higher energy X-rays have worse absolute energy resolution (∆E)
with more charges but better relative energy resolution (∆E/E) with better statistics.
The ground calibration verified the energy resolution of ∼180 eV at 5.9 keV (FWHM).
The CTI normally gets worse through the mission with in-orbit radiation damage. The
instrument and science operations teams will calibrate the energy resolution regularly withthe 55 Fe calibration sources and/or celestial source observations and updates the result in
the calibration database.
Background: Cosmic Ray (CR) particles or fluorescent X-rays from detector materials
excited by those particles interact with the CCD sensors and produce charge clouds in
CCD pixels. The grade selection method effectively excludes those events but selected(cleaned) data still include some of them called Non-Xray Background (NXB). The NXB
flux changes with cutoff rigidity in orbit, whose distribution varies with the 11-year Solar
cycle. (Figure 6.4 left, Nakajima et al., 2018). The amount of NXB depends on the exposed
duration of each pixel on a CCD and, therefore, the CCD rows (vertical pixel coordinates).
The instrument and science operations teams will study the NXB contribution to cleaned
data under various observatory environments and provide a tool that evaluates the NXB
spectra during individual observations. The sky background contribution is expected to be
similar to the Hitomi SXI at 5.6×10−6 counts s−1 arcmin−2 cm−2 in the 5−12 keV band
(Figure 6.4 right, Nakajima et al., 2018).

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Table 6.2: Available Window/Burst Mode Options

N

Image Size Time Res. Exp. time LTF
Pile-up(H) × (V)
(sec)
(sec)
(mCrab)Full Window No Burst
1
640×64043.96310.99
2.5
1/8 Window No Burst
8
640× 800.5
0.46310.93
21
1/8 Window 0.1s Burst 8
640× 800.5†
0.06200.12
160
N : Number of exposure in a unit time interval (4 sec).
Image Size: each CCD chip size.
LTF: Live Time Fraction.
Pile-up: On-axis source flux that suffers 10% photon pile-ups, equivalent to an
SXI count rate at 6.3 cts s−1 .
† Users can measure each photon’s arrival time information with ∼0.06 sec accuracy.

6.2

Observation Options for General Users

Users can choose the following optional operation modes for CCD ID=0 and 1 to observebright sources and/or improve time resolution (Table 6.2). These two sensors run in thesame mode as they share the CCD sequences, while the CCD ID=2&3 sensors always run
with the Full window mode.

CCD_ID=1
Resolve FOV
CCD_ID=0

CCD_ID=2

CCD_ID=3!I

Figure 6.5: Simulated SXI image of the Perseus cluster with the 1/8 window mode. The green boxes
show the approximate chip boundaries.

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Window option
The window option reads only a divided portion of the CCDs. The guest observer’s programsupports the 1/8 window option, which reads only 80 CCD rows for one exposure. This
operation reduces the readout time to 1/8, so the exposure cadence, and therefore the time
resolution, increase by 8 (see Figure 6.5).

Burst option
The burst option collects events obtained during a fraction of the nominal exposure time.
This mode reduces photon pileups of bright sources with short-frame exposures. However,it only takes a single exposure in an exposure cycle and trashes events outside of the
exposure interval. So, it does not increase the time resolution. The product data havetime stamps of the beginning of each exposure so that users can refine photon arrival times
for timing analysis. The guest observing program supports the 0.06 sec exposures (12%duty cycle) for the 1/8 window mode.

How to Choose the Observing Mode
Many point sources optimal for Resolve observations would exceed the Full window mode
pileup limit. Even some diffuse sources like Perseus cluster, Pup A and Cas A have brightknots that go beyond the pileup limit (see e.g., Figure 3 in Tsunemi et al., 2013). Which
Xtend mode should users choose?
The answer depends on the goal of their Xtend observation. If users want to collectphotons only from the aim point target as many as possible, they should choose the 1/8
window mode and add the burst option if the target is brighter than ∼21 mCrab. However, the pileup occurs only near the PSF core, so users may be able to utilize data fromthe extended PSF outskirts with little photon pileups using the Full window mode. The
HEASOFT tool, pileest, would help estimate pileup fraction from obtained event data.

6.3

Things to Be Considered for Bright Source Observations

6.3.1

Photon Pileup

Two or more X-ray photon events are indistinguishable when they fall in 3x3 CCD pixels
in a single frame exposure; such an incident is called photon pileup. Many pileup events
have two photons at different pixels, recognized as a multiple-pixel event and discarded,
while others are just counted as single X-ray events. Either way, the SXI counts two or
more pileup X-ray photons as one or zero, underestimating the photon flux. If detected as
single X-ray events, pileup events have an energy of two or more X-ray photons, making

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the spectrum harder. The PSF core has the highest count rates per CCD pixel and prone
to pileup. The PSF peak diminishes as the source flux increases and turns into a hole in asevere case.
The pileup fraction is the ratio of pileup events over all incoming X-ray events. Xtend
SXI reaches the 10% pileup limit at a count rate of 6.3 cts s−1 , corresponding to ∼4.5
×10−11 ergs cm−2 s−1 between 2−10 keV for the Crab nebula spectrum with a power-law
Γ at 2.13 (an IACHEC model2 ). Pileup tolerance depends on users’ science goals. For
example, users should minimize the pileup to detect a weak hard spectral tail, while they
may allow some degree of pileups to detect emission lines.

6.3.2

Out-of-Time Events

Out-of-time events are X-ray events detected during the CCD frame transfer — paralleltransfer of charges from the imaging area to the frame store area. One logical pixel transfer
takes only 57.6 µsec, so only bright sources appear as vertical streaks above the backgroundin images. Out-of-time events rarely suffer photon pileups as detected charges stay in a
fraction of time in a pixel, and so can be used for measuring non-pileup spectra of extremely
bright sources.
In the normal clocking mode, out-of-time events constitute a uniform brightness streak
and contribute only ∼0.93% of the total counts from a target. In the burst option, they
constitute a significant fraction of the source counts with a relatively long frame transfer
interval to the frame exposure, contributing 7.3% of the source count during the exposure.

6.3.3

Sacrificial Charge

Very bright sources create so many X-ray events that their charges also work as fillers ofcharge traps in the CCD chips (?). This effect improves the spectral resolution, but it
depends on the source brightness and is difficult to estimate. The instrument team does
not plan to implement the effect in the standard calibration.

6.4

SXI in Depth

The following sections describe the detail of the SXI instrument. They may not necessarily
be for proposal preparation, but they should help understand how the instrument works.
The SXI sensor (SXI-S) is placed at one of the two XMA’s focal plane — 5.6 m below
the mirror — on the XRISM ’s optical bench. The SXI instrument system consists offour components: a sensor body (SXI-S), pixel processing electronics (SXI-PE), digital
2

https://iachec.org

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electronics (SXI-DE), and a cooler driver (SXI-CD). The SXI-S (Figure 6.6) includes the
CCDs and their supporting systems, driver and video electronics, a single-stage Stirlingcooler, and calibration sources. The SXI-PE extracts X-ray event information from initial
image data. The SXI-DE packages X-ray event and associated housekeeping (HK) data
and delivers them to the Satellite Management Unit (SMU). The SXI-DE controls the
Stirling cooler and associated heater, keeping the CCD sensor’s temperature at −110 ◦ C(early operational phase), −120 ◦ C (late phase), or the desired temperature.
Table 6.3: SXI Characteristics

CCD Type
Number of CCDs
Number of segments per CCD
Imaging area
Pixel size
CCD gap
Depletion depth
Optical blocking layer
Contamination Blocking Filter (CBF)
CCD operating temperature
Calibration sources

6.4.1

Frame transfer, p-channel, Backside illuminated4
2
31 × 31 mm2 , 640 ×640 logical pixels(1280×1280 physical pixels)
24 µm × 24 µm (physical pixels)
1.2−1.6 mm (44−59′′ )200 µmAl(100 nm)/Al(100 nm) deposited on the CCD surface
Al(80nm)/Polyimide(200nm)/Al(40nm)−110 ◦ C/−120 ◦ C
two 55 Fe radioisotopes

X-ray Detection Mechanics and Operation

Hamamatsu Photonics K.K. fabricated twelve flight-model candidate CCD chips, and the
Xtend team selected four CCD flight-model chips among them (Yoneyama et al., 2021).
The SXI CCDs have a p-type channel and n-type substrate, in which positively charged
holes act as mobile charge carriers instead of electrons for earlier X-ray CCD cameras (SeeFigure 1 in Bebek et al. (2004)). The high resistivity of the n-type substrate produces athick depletion layer so that the CCDs are fully depleted to the back surface with a depthof ∼200 µm. For comparison, the Suzaku XIS backside illuminated CCD with a p-type Si
substrate was polished to 45 µm to ensure full depletion. The large SXI CCD depletion
depth combined with the backside illumination improves the quantum efficiency at both
soft and hard X-ray energies.
The larger depletion depth causes longer drift length and more extensive diffusion.
This effect strongly affects soft X-ray photons absorbed near the CCD surface: because
the CCDs are backside illuminated, their charges need to drift from near the surface to thebottom with the electrodes. X-ray event charges rarely fall in a single physical CCD pixel(24 µm × 24 µm) but multiple pixels. To simplify the event processing, the SXI nominallycombines charges in every set of 2×2 original pixels before the CCD readout and converts

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Figure 6.6: External View of SXI-S

them to a digital value. This 2×2 onboard binning format has ∼30% of X-ray events in asingle binned pixel, and the rest within 3×3 binned pixels. In the following, the ”physical”
pixel refers to the original 24 µm × 24 µm CCD pixel, while the ”logical” pixel does the
2×2 binned 48 µm × 48 µm, pixel. The SXI science data output in the 640 × 640 logicalpixel format for each CCD.
Each CCD chip has a rectangular shape with a square imaging area (IA) and frame
store (FS) area. The IA has 1280 × 1280 physical pixels, exposed to the sky through the
XMA for X-ray photon detection. The FS also has 1280 × 1280 pixels but is covered witha metal plate and used solely for temporary charge storage. In a normal data acquisition
sequence (Figure 6.7), the IA area detects X-ray photons in a frame exposure (3.96 sec in
the normal clocking, Full window mode), whose charges quickly drift to the buried channelof the corresponding pixels. After the exposure, the CCD clock driver changes the electrodevoltages (so-called clocking) to move the accumulated charges by one vertical pixel at one
cycle and repeat this cycle 1280 times to transfer all charges in IA to FS. This “paralleltransfer” process transfer all charges in the 1280 CCD columns in parallel. One pixel shift
takes 28.8 µsec, and a whole parallel transfer from IA to FS takes 36.9 msec. During eachparallel transfer, the SXI also operates charge injection for the subsequent frame exposure(Section 6.4.2 for the detail) so that the parallel transfer duration depends on the charge

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injection frequency and the clocking mode
After the parallel transfer, the CCD driver transfers charges in the bottom row horizontally to readouts with two serial registers. A serial register handles a CCD segment,
moving charges in the bottom row by one horizontal pixel at one clock cycle toward theactive readout node. An analog-to-digital converter measures the number of charges at
the readout and stores the value in memory. Once the serial register reads charges in thewhole pixels in the horizontal line, the driver vertically transfers all FS charges by one
pixel and then read charges moved to the bottom row. The driver continues this sequence
until reading all FS charges.

Figure 6.7: Schematic diagram of the charge transfer process on a CCD. (a) At the end of a frame
exposure, X-rays or charged particles produce charge clouds around the reacted pixels (red dots). (b) The
CCD clock driver transfers the charge clouds vertically to the FS region. It also injects artificial charges
to the top serial register every N vertical transfers (N =80 in the prelaunch default setup), which move
down to the IA with the subsequent vertical transfers (red lines). (c) Once all IA charges move to the FS
area, the CCD clock driver only transfers charges in the bottom FS row horizontally for readout and thenmoves all FS charges down by one pixel for another horizontal readout. It repeats this process until all FS
charges are read out. IA takes another exposure during this readout process. (d) Another charge transfer
process begins. The CI spacing and X-ray cloud sizes are not to scale.

6.4.2

Charge Transfer Inefficiency and Charge Injection Operation

Defects in the CCD silicon lattice can hold charges for a short time. If a defect, or socalled trap, is on an electrode, it can capture a portion of charges under transfer andkeep them until untrapped charges move to a downward pixel. This process, called chargetransfer inefficiency or CTI, occurs stochastically, producing a tail to the original charge
distribution and adding noise to the reconstructed event pulse height. The chance of trap
increases with the number of pixel transfers, so more charges tend to be ripped off fromevents detected further from the readout node (i.e., large RAWY coordinates). The SXIchips have relatively large CTIs in the ground testing, and the CTI is expected to progressin orbit as particles continuously bombard CCDs and break the silicon lattice (Kanemaru

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Figure 6.8: Results of a ground test of charge injection (CI) operation in September 2021 with an SXI-S
CCD engineering model. In this test, 5 keV charges were injected to the CCD chip by every 80 binned
rows. Left — Peak Pulse Height (PH) values of Mn Kα source events with/without (red/black) CI with the
ACTY row, equivalent to twice the number of parallel transfers. CI prevents PH (i.e. charge) reduction.
It produces a saw-tooth pattern because it works better on pixels following near a CI row. Right — 55 Fe
radioactive source spectra accumulated from all events on the CCD chip. CI recovers the line energy andthe energy resolution.

et al., 2019).
To ameliorate the CTI problem, the SXI employs the charge injection (CI) technique.
The Suzaku observatory’s X-ray CCD camera (XIS) employed this technique for the firsttime in orbit and substantially improved the energy resolution in the middle of the mission to near the initial performance (see, for example, Bautz et al., 2007; Uchiyama et al.,2009). This technique artificially fills a large number of charges every N CCD row, which
fill traps before X-ray events pass the traps during the vertical transfer (Figure 6.7). Figure 6.8 clearly shows that the calibration source spectra improves with the CI operation.
A downside is that the CI rows stay in the IA during frame exposures, so these lines, andprobably their leading and trailing rows, cannot detect X-rays. This limitation reduces theeffective area by 1/N or more.
Most trapped charges escape quickly and appear in a few trailing pixels. These charges
change the original 5×5 pixel patterns, which is problematic for the grade event extraction
classification. This charge migration can be statistically estimated and restored, and so thecharge trail correction algorithm runs during the pipeline processing to fix the problem. A
small amount of charges reappear outside the 5×5 pixels, causing an artificial pulse height
deficit. The CTI correction algorithm in the pipeline processing recovers the expected loss
based on the ground and onboard calibration study (Yoneyama et al., 2020).

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59

Hot Pixels

Defects in the CCD silicon lattice can also produce charge currents without X-ray orparticle events, which appear as high pulse height pixels called hot pixels. Some hot pixels
are permanent, always showing pulse heights above the defined hot pixel threshold. SXI-PE
finds such pixels when taking dark frames onboard and flags them as hot pixels. The otherhot pixels, specifically flickering pixels, show high pulse heights randomly or occasionally.
They may not appear as hot pixels in dark frames, so ground data analysis finds themthrough statistical analysis.
Strayed optical lights can produce pseudo-hot pixels under pinholes of the optical blocking layer. If the known pinhole pixels have a significant event rate increase in orbit, theinstrument and science operations teams will obtain frame images and investigate thecause. The teams do not expect optical light to leak into the SXI camera with the current
observatory and instrument designs.

6.4.4

Area Discrimination

If a CCD constantly outputs enormous numbers of events from a very bright source or hot
pixels, the telemetry may be unable to handle all event data and lose some onboard. If it
occurs, the science operations team may choose CCD areas to downlink to avoid telemetry
saturation. This area discrimination setting is not a selectable option for general observers.

6.4.5

Supporting Components

Optical Blocking Layer and Contamination Blocking Filter
X-ray CCDs are also sensitive to optical and infrared lights, which can work as noise for Xray observations. To block the lights, the SXI CCDs coat the surface with Optical Blocking
Layer (OBL), two Al layers 100 nm thick each. OBL blocks optical lights to infrared, while
it transmits X-ray emission above ∼0.4 keV (Figure 6.3). The OBL has a small numberof pinholes of 0.1−0.3%, which might produce high pulse height pixels with stray opticallights though the instrument team does not expect such lights come into the camera body.
Pixels under the pinholes are flagged and may not be used for science data analysis if they
produce significant noise.
The SXI CCD, with an operating temperature of −110/−120 ◦ C, is significantly colder
than the other satellite components. Adhesive outgases from the satellite body tend to
stick to such cold surfaces permanently. The Suzaku and Chandra observatories had suchbuildups on the optical blocking filters of their CCD instruments, which significantly reduced the soft X-ray sensitivities in-orbit. The Xtend SXI includes CBF at the hood top,which is composed of polyimide with a 200 nm thickness sandwiched between aluminumProposers’ Observatory Guide

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with thicknesses of 80 nm and 40 nm. It blocks outgases from getting inside the SXI bodywhile it is warm enough at about +25 ◦ C at 50 cm above the CCD chips not to build up
contamination on its surface. The CBF also eliminates a vacuum-tight chamber and dooropening mechanism (Mori et al., 2022). Figure 6.3 shows the combined quantum efficiencyof the CCDs, OBL, and CBF.
Calibration Sources
The SXI has two 55 Fe radioactive sources installed into the bonnet for in-flight calibration(Figure 6.6). The sources emit Mn Kα (5.9 keV) and Kβ (6.5 keV) radioactive decay lines
to two circular spots on the CCD detector plane (Figure 6.2). The instrument and science
operations teams will collect in-orbit data of these lines and measure each CCD chip’s gainand energy resolution from their line centroids and widths. The team regularly updates
the result in the calibration database so that users can utilize the latest calibration by
reprocessing their datasets.
The holder of each calibration source collimates direct radiation to ≲4.8 mm (∼3′ )from an adjoining corner of two CCDs (Figure 6.2). The ground calibration cannot fully
reproduce the in-orbit condition, so that the instrument team will measure the irradiatedareas in orbit.
Mn MainKαKβ
Mn
EscapeKαAlKβFigure 6.9: 55Fe calibration source spectrum, obtained during the thermal vacuum test in August 2022.

Cooling Systems
A mechanical Stirling cooler removes heat from the SXI’s focal plane at a constant rate. The
SXI-DE monitors the CCD temperature every second and powers the heater on the backsideof the cold plate to regulate the focal plane temperature at −110/−120 ◦ C. The CCD
temperature was stable on ground EM testing to within 0.1 ◦ C. Such a low-temperaturefluctuation helps stabilize the CCD performance and ease the calibration effort if realizedProposers’ Observatory Guide

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in orbit.

6.4.6

Change from the Hitomi SXIThe Xtend SXI is almost identical to the Hitomi SXI (Tanaka et al., 2018; Nakajima etal., 2018), but there are design changes to fix the problems found in the Hitomi SXI.
Adding a Notch Structure on the CCD Electrodes: The Hitomi SXI has arelatively large charge transfer inefficiency (CTI), degrading spectral energy resolution (e.g.,
Kanemaru et al., 2020). The Xtend SXI CCDs add a notch implant to each channel, whose
potential well produces a narrow charge transfer path, reducing the chance of encounteringcharge traps. With the notch, the Xtend SXI has a factor of 3 smaller CTI than the Hitomi
SXI. The Xtend SXI also does not show spatial variation in CTI seen in the Hitomi SXI.
Coating Additional Al Layer to the Optical Blocking Layer: The optical blocking layer of the Hitomi SXI had multiple pinholes with high optical/IR light transmission.
When scattered optical/IR lights leaked into the Hitomi SXI body, these pinholes produced
large pulse height pixels. The Xtend SXI adds a 100 nm thick Al coating to the original
100 nm Al coating on the CCD surface, reducing the number ratio of pinhole pixels from
4.1% to 0.2% (Uchida et al., 2020). This coating, in turn, slightly degrades soft X-ray
sensitivity.
Placing an Aluminum Sheet on the Wiring Area: The bonding sheet between the
CCD and Si base is transparent to optical lights. When scattered optical/IR lights leaked
into the Hitomi SXI body, the lights sneaked in through the bonding sheet, producing
pulse height increases of pixels near the CCD edges. The Xtend SXI makes Al coats on
the layer of the backside electrode (under the passivation) to block the optical light.
No Openings on the Optical Bench: The scattered lights mentioned above originated from the two openings in the fixed optical bench for the Hitomi hard X-ray instrument (HXI). XRISM has no openings on the optical bench without the HXI and so expects
no incoming scattered light inside the satellite and SXI body.
Shifting the Aim Point: The Xtend SXI’s aim point is placed at (8 mm, 8 mm) fromthe two active pixel edges3 . These values are larger than the Hitomi SXI values (6.5 mm,
6.5 mm), intended to minimize the portion of the main target’s PSF outskirt falling in thechip gap.

3

The precise values are measured with onboard calibration.

Chapter 7

Observations of Extended Sources
In this chapter, we present the characteristics of XRISM instruments that are relevantto the observations of extended sources. In particular, we discuss the possibilities andlimitations of spatially-resolved spectral analysis given the imaging performance of XMA
and the smaller FOV of Resolve. Careful simulations are often necessary to assess, at areasonable degree of confidence, whether a particular spatial-spectral analysis objective isachievable.

7.1

Challenges of analyzing extended sources

The non-dispersive high-resolution spectroscopy offered by the Resolve instrument onboard XRISM allows one to observe extended sources (e.g., clusters of galaxies, supernova
remnants) with an unprecedented spectral resolving power (Chapter 5). Similarly, theses
sources can also be imaged and investigated spectrally (though at much lower spectral
resolution) using Xtend (Chapter 6). However, the limited spatial resolution permitted by
the observatory (as well as the limited number of pixels on the Resolve detector) requires
extra caution from observers of these systems.
The two main limiting properties of Resolve are:
• The Resolve pixels are 0.5′ × 0.5′ , i.e. substantially smaller than the PSF (∼1.3
arcmin HPD);• The Resolve field of view is 3′ × 3′ (i.e. of size comparable to the PSF).
The Xtend instrument has a larger field of view and, therefore, is not affected by thelatter issue. The PSF, however, is caused by the XMA (Chapter 4) and thus affects Resolve
and Xtend in the same way.

62

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As a consequence, in both instruments photons from specific regions of the sky are expected to spread over the detector and “contaminate” (more specifically, mix with photons
from) other regions. Given that this process occurs at all energies, this mixing will affectthe source not only spatially but also spectrally. In the following, we refer to this process
as “spatial-spectral mixing” (SSM).
Concretely, and depending on the extent, properties and environment of the source,
Resolve can be affected by two types of SSM:
• Internal SSM: when the emission within the field of view mixes across pixels andcontaminates other regions within the same field of view. For instance, a supernova
remnant hosting clumps and/or hot spots with specific spectral features – e.g., brightmetal lines – or a complex star-forming region with several spatially-resolved stars.
This signal will mix through the entire region covered by Resolve, and one will notice,e.g., an enhancement of the same metal lines in other pixels as well.
• External SSM: when sources outside the detector contaminate the detector region
itself. One typical example of external SSM would be a pointing of the outskirts ofa relaxed galaxy cluster. Another example is the case of a bright AGN located a
few arcmin away in projection from any extended source of interest. In the former(latter) case, a fraction of photons from the bright cool core (AGN) leaks into thedetector and contaminates the observed outskirts (extended source).
The external and internal SSM are illustrated in Figure 7.1 for the simple case ofphotons from a single point-like source spreading over the detector. As seen on the figure,due to mirror support structures, imperfect alignments of telescope components, mirror
foil figure errors, and non-ideal foil surfaces, the PSF has a complex shape and is not
azimuthally symmetric. Particularly in the case of a complex extended structure (e.g.,with resolved point sources, clumps and/or hot spots), the PSF effects cannot be seen as asimple Gaussian-like smoothing of an ideal image and should be thus addressed carefully
and accurately. The next two figures (Figs. 7.2 and 7.3) focus on a concrete (thoughvery simplistic) case of a circular extended source, half of which is redshifted (V1 = +235
km s−1 ) while the other half is blueshifted (V2 = −235 km s−1 ). At first approximation,
this is close to what is expected for galaxy clusters with rotating bulk motions of their
hot haloes. The spectral effects of the PSF mixing of that specific source are shown in
Figure 7.3. Quite interestingly, a mixing of the two components with equal fluxes (e.g.,
if one looks at central region covering the two components simultaneously) produces adouble-peaked line profile, much easier to identify and interpret than the asymmetricalprofile seen when one component dominates over the other (e.g., if one looks at a region
with one component – though inevitably affected by SSM still).
Generally speaking, the SSM issues described above are somewhat less impactful for
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Figure 7.1: Shape of the PSF and its effective impact on Resolve. Top panels: case of a point-like source
placed at the center of the field of view (left), translating into significant internal SSM on the Resolve
pixels (right). Bottom panels: case of a point-like source placed directly outside of the field of view (left),translating into significant external SSM on the Resolve pixels (right). “BL-CCD” stands for “Beam-LineCCD”.

of view. Except in specific cases, a moderately extended source placed on axis of the(large) Xtend field of view will thus make the external SSM less important in comparison,

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Figure 7.2: Simplistic case of spatial mixing for a circular extended source whose left and right halfs are
redshifted and blueshifted, respectively (uppermost panel). The two intermediate panels show, from that
region (grey dotted circle), a distribution of all the incoming (redshifted or blueshifted) photons eventuallyreaching the Resolve detector (with yellow square and dashed lines lines indicating the Resolve detector
boundaries).The XRISM PSF effectively mixes photons from each part into the other, resulting in 15-17%of cross-contamination (two bottom panels, showing the fractional number of photons).

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and it will be easier to quantify the internal effects of the PSF mixing within the source.
Moreover, the moderate spectral resolution of Xtend means that there will inevitably beless subtle spectral features to be affected by SSM. This, however, does not mean that SSM
effects can be neglected when analyzing extended sources with Xtend. In fact, variations
from e.g., bright emission lines (and/or the continuum) are still important to take into
account even at CCD resolution.
The “advantage” of the SSM from Xtend is that the issue is not new. In fact, such
effects have affected extended sources in very similar ways in previous missions such as
ASCA or Suzaku. For this problem, we thus refer the proposer to the literature whichabound with many concrete cases (for the case of galaxy clusters, see e.g., Markevitch
et al. , 1996; Bautz et al., 2009; Bulbul et al., 2016). In contrast, the above-mentioned
characteristics of Resolve make (internal and external) SSM issues more complicated andrelatively unexplored (see however the case of the Perseus cluster observed by Hitomi ; e.g.,
Hitomi Collaboration , 2018). The rest of this chapter will therefore focus primarily on
Resolve.

7.2

Facing the challenge: methods and prospects

As mentioned earlier, the problem of SSM is not completely new, as it has already beenencountered to some extent in other X-ray missions – in particular those with a large PSF,
such as Suzaku and NuSTAR. The spectral analysis of X-ray extended sources in thosecases is often performed as follows:
1. Subdivide the extended source in various spatial regions of interest (e.g., where spectral features and physical properties may vary);
2. Extract, for each spatial region, a spectrum (Sj ) as well as an associated ARF (Aj ).
In most cases, the same RMF (Rj ) can be assumed for all regions as it weakly varies
across the detector;
3. Fit each region with a linear combination of models, where each model Mi is associated with each spectrum Sj as follows:

S1 = A1 R1 (P1→1 M1 + P2→1 M2 + P3→1 M3 + . . . )(7.1)

S2 = A2 R2 (P1→2 M1 + P2→2 M2 + P3→2 M3 + . . . )(7.2)

S3 = A3 R3 (P1→3 M1 + P2→3 M2 + P3→3 M3 + . . . )

(7.3)

...

(7.4)

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Under a more generic form, one has:X
Sj = Aj Rj (Pi→j Mi ),

(7.5)

i

where Pi→j are coefficients corresponding to the relative contribution of the model
Mi in the spectrum Sj . In an ideal case with no contamination at all, the Pi→j matrix
is in fact diagonal. Even in more realistic cases, many Pi→j terms may be negligible(particularly if regions i and j are more than a few arcmin away), however the SSM
will inevitably make several other Pi→j values positive.
To estimate numerically each value from the Pi→j matrix, a tempting approach would
be to associate them directly with the normalization parameter of each model Mi in simultaneous fits of all spectra Sj . Such normalization parameters (Ni→j ) would then be left freein the fits. This method, however, is strongly discouraged for a number of reasons. First,
thawing all these normalizations considerably increases the number of free parameters in
the fits, with a high risk of reaching strong degeneracies between the latter and/or local
minima in the fits. Also, the best-fit values of the free Ni→j coefficients may not reflect
the true mixing rate in the considered region (as such normalizations could be biased byother unaccounted uncertainties). In fact, free normalizations between the matrix elements
uncouple the problem from the actual telescope and detector, making a “black box” of thewhole SSM issue, and will thus likely give artificial results. Such approach should be thus
avoided by all means, particularly in the case of Resolve observations.
A much more accurate way to proceed is to determine these coefficients directly by
calculating (i.e. via ray-tracing) the contribution of each component from region i into aconsidered region j. This can be done from the XRISM data reduction software by theroutine xaarfgen (Figure 7.4) which, via its ray-tracing subroutine xraytrace, generates
an ARF that can be applied directly on a given model Mi . The obtained effective area has
been specifically scaled by the routine to account for the exact contribution from region
i into region j, meaning that the normalization parameter of each model Mi remains untouched (thus, it conserves its initial physical interpretation). The subroutine xraytrace
is able generate such ARFs from a geometrically simple input model (e.g., point source, βmodel) or even an input image; e.g,. taken from Chandra/ACIS, XMM-Newton/EPIC, or
simply from a simulation1 . Adopting this method, the Pi→j values have their information
directly encoded in a grid of ARFs generated specifically for this purpose (Ai→j ), and the1

Despite this being the best feasible approach by far at the moment, it is good to keep in mind thatgenerating an ARF in a local region from an input image has its own limitations. In particular, suchan image stands for one specific energy and does not account for the energy-dependency of the extended
source. Such systematics can be assessed by reproducing the same exercise with input images obtained
from different energy bands. In the case of the Perseus cluster, for instance, the energy dependency of thismethod using Hitomi SXS data was found to be negligible (Hitomi Collaboration 2018, PASJ, 70, 9).

Proposers’ Observatory Guide

68

above system of equation becomes then:S1 = R1 (A1→1 M1 + A2→1 M2 + A3→1 M3 + . . . )(7.6)

S2 = R2 (A1→2 M1 + A2→2 M2 + A3→2 M3 + . . . )(7.7)

S3 = R3 (A1→3 M1 + A2→3 M2 + A3→3 M3 + . . . )

(7.8)

...

(7.9)

or, under its generic form,

Sj = Rj (X

Ai→j Mi ).

(7.10)i

In the case of Resolve, it becomes then tempting to follow the same method with,for bright and very extended objects, up to 35 regions (each corresponding to one active
pixel on the detector). However, the wealth of spectral information delivered by XRISM
associated with the complexity of the ray-tracing process used for the mission requires
computing resources that may become considerably heavy. In fact, the computing time
requested for ray-tracing in the “fast-mode” (i.e. with more limited off-axis accuracy) andextracting a single ARF can easily exceed the hour, making the above approach extremely
time-consuming (i.e. on the order of days or even weeks for standard single machines).
Besides the use of distributed/pseudo parallel computing (if available at the user’s hostinstitute), a library of pre-computed PSFs (as a function of the energy, off-axis angle,
azimuthal angle) will be made available for Resolve after the launch of the mission (likely
ready for AO-2). The sum of contributions will then generate a series of fine-grid effective
areas for the spectral regions considered. Another alternative to reduce computing times
is to focus on a small energy range2 (if permitted by the science goal) and/or reduce thenumber of ray-tracing photons. Also, since the SSM is due to the telescope, the ray-tracing
code xrtraytrace as part of the XRISM data analysis software can be used to make SSM
estimates by running standalone for single energies. For proposals and/or actual analysisof real data, one can still use the above approximate methods to make an initial assessmentof which cross-mixing elements are negligible (e.g,. two regions that are distant from eachother). Doing so, the run time for creating the full ARFs is not wasted on ARFs that canbe safely omitted.
2

A single energy is not possible with the current ARF generator, but a narrow energy band is.

Proposers’ Observatory Guide

7.3

69

Consequences and general considerations

As we have seen in the previous two sections, the analysis of extended sources with XRISM
is possible, but challenging. To help proposers evaluate to what extent their science goal
is feasible, we provide in this section a list of items and considerations that are importantto keep in mind at all stages of the observation – from the proposal to the analysis.
We stress that this list is purely indicative, non-exhaustive, and does not constitute a
formal checklist for the technical feasibility of any future proposal.
1. It is very likely that the proposers’ extended source will be affected by SSM. To
estimate how the latter will impact the analysis of the former (and the proposedscience goals), it is important to take into account all information on what is known(or unknown) about the spatial and spectral structure of the source. The presenceof notable features (e.g., hot spots, blobs, bright nearby AGN or centrally peaked
emission), each with their own spectral behavior, will inevitably affect the spectrumof the spatial region(s) to be investigated. Determining how strongly this will affectyour science goals and what strategy to use when accounting for this effect will beessential for conducting successful observations.
2. If known to some extent, it is strongly advised to plan a coherent mapping/tiling of
the source; i.e. to define spatial regions where pixels are thought to have a (more orless) similar spectral behavior. This also means that, in some cases, the observation
should rather be aligned for an optimal match of pixels with spatial features of thesource (e.g., to better capture cavities in clusters or shocks in SNRs), rather thansimply match the center of the source with the on-axis position of the detector.
3. As the science goal might depend on which region of the extended source (and withwhich roll angle) the telescope should point at (e.g. if one wants to minimize the SSM
in one particular region of interest), it is also important to check the target visibility.
This can be done using the visibility checker tool and optimizing the observational
position based on the position angle in the observation period.
4. Generating ARFs can be computationally expensive. This is especially true in thecase of extended sources, as accounting for the SSM often requires to generate alarge number of ARFs. The proposers should keep in mind that several methods
exist, some may be more appropriate depending on the proposers’ needs and available
computing resources (e.g., computing precise ARFs directly, using fastmode, using
PSFLib, using an alternative or hybrid method, etc.).
5. A consequence of the previous item is that, although tempting, using Resolve formulti-pixel array spectroscopy is challenging: given that any given pixel is likely
affected by the emission from the 34 others, this approach requires the simultaneousProposers’ Observatory Guide70

fitting of at least 35x35 = 1225 spectral components3 . If only one pixel is of particularinterest, however, the problem can be considerably simplified by grouping all theothers in one common region.
6. The SSM is commonly (and correctly) assumed to happen within the field of view.
The external SSM, however, should not be ignored, as its contamination to the investigated sky region(s) can be non-negligible in some cases.
7. Quite counterintuitively, substantial SSM is not necessarily “bad” for the analysis.
For instance, a mixing of two components in equal proportions is likely to give rise tospectral features that are easier to spot (and control) than a fainter contamination
whose subtle features may lead to significant biases. Specifically, two equally brightregions of different velocities could result in a double-peaked line, much easier todisentangle than a 90-10% mixing resulting in slight (yet significant) line asymmetry.
8. Once the data are retrieved, spectral clues of SSM should rather be investigated byeye before performing a blind fit. Fitting residuals might also provide some clues
for mixing (e.g., finding a bump where no line is expected). Similarly, even at the
proposal stage it is encouraged to explore the feasibility well beyond a simple setup
of fakeit and fitting commands in Xspec.
9. The proposers should keep in mind all other systematic uncertainties that may affectthe analysis of their extended source. Effective areas, for instance, have their own uncertainties. Those can be of various origins (e.g., calibration, statistical fluctuations
in ray-tracing estimates, or even uncertainties related to software algorithms and assumptions). Similarly, other particular properties of the source might, in some cases,
further complicate the analysis (e.g., pileup effects – see Chapter 8, uncertainties inatomic databases impacting spectral features relevant to the science goals, etc.).

3

Taking the external SSM into account, each region requires at least one additional off-detector region.
This means that the above number of components fitted simultaneously further increases to 36 × 35 = 1260.

Proposers’ Observatory Guide

71

Figure 7.3: Spectral mixing effects on the Fe-K line at 6.4 keV as seen by Resolve, assuming the simplistic
case presented in Figure 7.2 (V1 = -235 km s−1 ; V2 = +235 km s−1 ). The top four panels assume a line
width of σ = 1 eV (FWHM = 110 km s−1 ), separating cases for contamination of equal fluxes (left panels)
vs. cross-contamination of 16% representing the regions shown in Figure 7.2 (right panels), as well as
theoretical models (upper panels) vs. more realistic profiles convolved by the Resolve RMF (lower panels).
The bottom four panels show the same effect for a larger line dispersion width (σ = 4 eV; FWHM = 440
km s−1 ).

Proposers’ Observatory Guide

72

Figure 7.4: Illustration of the tasks performed by the routine xaarfgen—and the utility of the latter incomputing SSM coefficients for the analysis of extended sources.

Chapter 8

Resolve Observations of BrightSources
Bright sources, here defined as those that produce ≳ 100 ph s−1 in the Resolve bandpass,
are important XRISM targets because the high count rate enables precise measurementsof variability in emission and absorption lines, as well as detecting weak features at highsignificance. The X-ray flux that produces a count rate exceeding 100 ph s−1 depends
on the spectrum, with softer spectra producing greater count rates at a fixed energy,
but as a rule of thumb this chapter applies to source fluxes exceeding 0.3 − 10 keV FX >5 × 10−9 erg s−1 cm−2 . This flux is equivalent to about 500 mCrab, where the Crab Nebula
X-ray flux is defined as 2 − 10 keV FX = 2.4 × 10−8 erg s−1 cm−2 (Willingale et al., 2001).
Almost all persistent “bright” sources are well-known Galactic compact objects(see https://heasarc.gsfc.nasa.gov/docs/heasarc/headates/brightest.html).
Above this flux, the number of source events (especially high-resolution events) nolonger depends linearly on the source flux and exposure time. High count rates can saturate pixels with low-resolution events, saturate the onboard event processor entirely, anddegrade the energy resolution due to electrical cross-talk (Mizumoto et al. , 2022). The net
impact is fewer clean, high-resolution events than expected for the exposure time. In thischapter, we briefly introduce the challenges posed by very bright sources, then describe
observing strategies to maximize the high-resolution event rate, and conclude with generaladvice on proposing and analyzing such sources. Since the Crab Nebula is the only current
extended source with a count rate that falls in this regime, observing strategies assume
point-like sources.
Any source that can produce ≳ 100 counts s−1 in Resolve will be extremely piledup in Xtend, for which a “bright” source is about thirty times dimmer (0.3 − 10 keV
FX ∼ 5 × 10−11 erg s−1 cm−2 ). Chapter 6 describes the impact of high count rates
on event detection (pile-up and out-of-time events) and the observing modes available to73Proposers’ Observatory Guide

74

mitigate issues. For almost any bright source, the optimal Xtend observing mode is awindowed burst mode.

8.1

Challenge: High-Resolution Event Rates

Observers should try to maximize the high-resolution yield, which we denote as Hp+Mp
following the convention of Section 5.3.3 (high-resolution and primary medium-resolution
events). Section 5.3 describes how events are detected. Section 5.3.3 describes how the
Hp+Mp yield depends on total incident count rate in a single pixel based on the characteristic time delay between photons: Hp+Mp increases with total rate up to a maximum
and then decreases as the average time delay between arriving photons becomes shorter
than the length of the high-resolution optimal filtering template. However, as almost all“bright” sources are point-like, we must also consider the illumination pattern of the Resolve array by the PSF. The total array Hp+Mp rate, Sarr,Hp+Mp , depends on the totalsource flux, the shape of the PSF and pointing center relative to the center of the Resolve
array, and the source spectrum. Sarr,Hp+Mp reaches a maximum of ≈100 counts s−1 foran on-axis point source with a Hitomi-like PSF (half-power diameter of 1.2′ ) at a total
incident rate of 1000 counts s−1 on the array.

8.2

Challenge: Event Loss

The Pulse Shape Processor (PSP) cannot process 1000 count s−1 ; its limit is about 200
count s−1 for an on-axis source, assuming that all events processed by the PSP are original
source counts (i.e., background events further reduce this). Events beyond this limit are
discarded and therefore “lost.” The PSP processes events by drawing the next event froma list provided from the event-finding board (a field-programmable gate array, or FPGA).
The packet of information from this list includes the addresses of the stored pulse data,so the PSP retrieves the pulse from the buffer and cross-correlates it with a template as
described above. The processing time depends on the template length. When processing
is complete, the event data is discarded and the energy, arrival time, etc. are sent to thedata recorder. If the data buffer fills up before the PSP has processed all waiting events,
the entire buffer is cleared and begins to fill up from zero. Hence, at high count rates the
PSP can be overwhelmed and events can be lost (but see below).
The PSP houses four CPUs, one for each quadrant of the 36-pixel Resolve array. EachCPU runs an event-processing application for each of 9 pixels so that more than one event
can be processed simultaneously. This means that event loss in one quadrant does notnecessarily imply event loss in other quadrants, which we will use to our advantage in
observing bright sources. When an event from a given pixel is processed, the default
behavior of each CPU is to put that pixel in the back of the line and start processingProposers’ Observatory Guide

75

events from the next available pixel. However, because different event grades take different
processing times, a pixel with many low-resolution events can get a higher fraction of its
event buffer processed than one with mostly high-resolution events. Averaged over thearray, high-resolution events can be preferentially lost.
Resolve does not “pile up” in the same way as CCD imaging spectrometers, but ananalogous effect can occur for two events that arrive at almost exactly the same time.
When two events are separated by less than a few ms, the secondary event-finding process
in the PSP can fail and treat the two events as a single event, thereby misreporting theenergy. Several percent of the events can be affected (Mizumoto et al. , 2022), with aconcurrent loss of several percent of the exposure time after screening.
Finally, it is important to note that the event-finding algorithm on the FPGA thatcreates the lists and fills data buffers for the PSP sends information about the numberof lost events to the data recorder, so not all information is lost when there is a buffer
overflow. The FPGA will not be overwhelmed for almost any conceivable X-ray flux andthis information can be used to reconstruct source fluxes (since whether an event is lost or
not, or whether it is high resolution or not, does not depend on photon energy). However,
without the search for secondary events performed by the PSP, the FPGA count rate is alower limit.

8.3

Challenge: Minimizing Cross-talk

In addition to event grading and event loss, a third challenge in the high-count-rate regime
is electrical cross-talk (Section 5.3.6). Voltage (“parent”) pulses in the wire from one pixel
induce “child” pulses in neighboring wires (thermal cross-talk is an order of magnitude
weaker and unimportant). Child pulses have an amplitude 0.6% that of the parent for
immediately neighboring wires and <0.1% for next-nearest wires. Due to the wiring map of
the array, child pulses are not necessarily induced in physically adjacent pixels (Figure 8.1).
Most pixels produce two child pulses per parent, and a few produce only one.
Because of their small amplitude, child pulses are rarely, if ever, detected as events intheir own right. However, child pulses make the signal noisier and can lead to significant
errors in the energy measured by optimal filtering if they occur at the same time as areal event. Such unhappy coincidences increase with count rate, so cross-talk can degrade
the energy resolution of the Hp+Mp spectrum for bright sources. This degradation varies
from pixel to pixel, so the total-array degradation is a weighted average. Under ordinary
circumstances, cross-talk contamination can be eliminated by ignoring events in electricallyneighboring pixels that occur within a few ms of each other. However, for very brightsources the fraction of events that occur within a few ms of each other may be quite high.
Strict filtering can eliminate many Hp+Mp events, making the cure worse than the disease.
The impact of cross-talk on scientific analysis depends on the spectrum and the goal.

Proposers’ Observatory Guide76

A hard spectrum at a very high count rate will produce child pulses with larger average
amplitudes than a soft spectrum at the same rate, leading to larger errors on the energy.
Likewise, if a line of interest falls at 1 keV the fractional error in the velocity resolution
will be worse than at 6 keV. The effect can be calculated exactly as a function of countrate for a monochromatic source, but for a more physically motivated spectrum we must
calculate it statistically. Figure 8.2 shows the impact of cross-talk for a 4 Crab (2-10 keV
FX ∼ 10−7 erg s−1 cm−2 ), on-axis point source with a Crab-like spectrum with no filters.
The left panel shows the broadening and centroid shift resulting from increased crosstalk contamination for an example Fe Kα line. The right panel shows the average energy
resolution (starting from a fiducial 5 eV resolution) in each pixel for this source (note thatthe most strongly affected pixels are at the array edge but are electrical neighbors of thecentral four). More examples of the impact of cross-talk on energy resolution are given forsimulated spectra in Mizumoto et al. (2022).
Finally, it is worth noting that in bursty sources (i.e., those with sharp variations within
a single observation) the spectral resolution may degrade during a particularly bright burst.
Observers interpreting line shifts and broadening during a burst must account for crosstalk.
Note: The estimates of cross-talk presented here are based on a representative pair of
pixels from the Hitomi Soft X-ray Spectrometer. Cross-talk impact can be non-uniform, as
it depends on the size of the parent pulse in volts per eV, the strength of the coupling, theinteraction of the cross-talk pulse shape with the unique digital filter of the parent pixel,
and the gain scale of the parent pixel. As noted above, the incident spectrum also makesa difference.

8.4

Observing Strategies: Filters and Off-Axis Pointing

Degradation of energy resolution due to electrical cross-talk becomes important above
0.3-10 keV FX ≳ 5 × 10−9 erg s−1 cm−2 . Event loss starts to become important above
FX ≳ 2 × 10−8 erg s−1 cm−2 . Above these thresholds, maximizing the “clean” Hp+Mp
rate (i.e., from pixels where the average cross-talk degradation is no more than 2% thenative energy resolution) is achieved by reducing the incident count rate from the default,
on-axis value. This can be achieved by using a filter to reduce the effective area (see curves
in Figure 8.3) and/or pointing off-axis. The Be filter removes most of the photons below
2 keV and is not useful in the current configuration, while the neutral-density (ND) filter
reduces the count rate at every energy by a factor of about four.
Pointing off-axis causes more focused X-rays to fall off the chip and can also be used
to concentrate low-resolution events and their attendant cross-talk signals in one sacrificial
quadrant, thereby maximizing the clean Hp+Mp rate in the other quadrants. Figure 8.4
shows the results for a Crab-like point source located at the quadrant opposite the calibration pixel. The upper right quadrant is saturated by low resolution events and contributes

Proposers’ Observatory Guide

77

24

26

34

32

30

21

22

25

33

31

29

19

20

18

35

28

27

9

10

17

0

2

1

11

13

15

7

4

3

14

16

8

6

5

12

Pulse
Normalized Pulse Height23

Cross-talk Pulse(shifted derivative,
0.6% pulse height)

0

20

40

60

Time (ms)Figure 8.1: Electrical cross-talk “child” pulses occur in electrically adjacent pixels, so the cross-talk
distribution does not follow the shape of the PSF. Left: Resolve pixel map showing electrically neighboring
pixels. A parent pulse in pixel 18 (blue) will produce a child pulse in pixel 19 (light blue), but not inpixel 17 because it is in a different and electrically isolated quadrant. Likewise, a parent pulse in pixel 31(orange) will produce child pulses in both pixels 30 and 32. Right: Cross-talk pulses (blue) have a shapebased on the time derivative of the parent pulse (red) and a much smaller signal, with a peak amplitudeof 0.6% of the pulse height. The child pulses are superimposed on the data stream in their pixels; whena child pulse in pixel 19, induced by a parent pulse in pixel 18, occurs around the same time as a parent
pulse in pixel 19, the pulse shape is distorted. This leads to an erroneous energy measurement.

about 10% of the Hp+Mp events, which also have average energy resolutions degradedby 0.1-2.5 eV. However, the other three quadrants produce about 55 counts s−1 with anenergy resolution within 0.1 eV of the fiducial 5 eV value. The count rates in the otherquadrants are also low enough that a crosstalk exclusion filter (±25 ms) can be appliedwith little or no impact on the number of Hp+Mp events, but such filtering should not
be necessary. Note: The structure of the PSF wings, based on a ray trace incorporating
ground measurements, is not yet fully known, especially as a function of off-axis angle.
Figure 8.4 thus does not have very high fidelity and is shown as an example. Observers
should also bear in mind the effects of roll angle.
Table 8.1 illustrates how the optimal configuration (in bold) changes with increasing
count rate for a sample spectrum from a point source. The Be filter is omitted as makes
hardly any difference over the gate valve itself (a much thicker Be filter). For sourcesbrighter than about 3 Crab, the ND filter becomes optimal; at 10 Crab off-axis pointing
also becomes important.
While this table is for the purpose of illustration (adopting a modestly absorbed Γ = 2
power law spectrum), it reveals the basic strategy: first try a filter, then point off axis

Proposers’ Observatory Guide

Normalized Flux

1.00

Fe K⍺1

78

Cross-talk Rate
0
1
3
10
30
50

0.10

0.01
6.375

6.385

6.395

6.405

Energy (keV)

6.415

6.425

Energy Resolution (eV)Figure 8.2: Left: Increasing contamination by cross-talk child pulses both broadens lines and shifts theircentroids. At high count rates, this can significantly alter the energy resolution in individual pixels, as
shown here for an example Fe Kα line. Right: As an example, this Resolve pixel map shows the energy
resolution (i.e., broadening) in each pixel when starting from ∆E = 5 eV and observing a source four times
brighter than the Crab Nebula at the array center. The four exterior pixels with the worst resolution areelectrically adjacent to the central four pixels. Note that the map is not weighted by the event rate or
grade: despite the better energy resolution in the central pixels, there are hardly any Hp+Mp events fromthose pixels.

towards the corner of one quadrant. We do not recommend using the quadrant with thedummy pixel, as the CPU for that quadrant must process the 55 Fe calibration source photons that are necessary for gain tracking. We have summarized general recommendations
as a function of source brightness in Table 8.2. Here we offer a few additional commentsfor proposers:
First, proposers should use the effective area curves (encoded in the ARFs) to estimate
the count rates for different filters (e.g., via WebPIMMS: https://heasarc.gsfc.nasa.
gov/cgi-bin/Tools/w3pimms/w3pimms.pl). On-axis, point-source ARFs are sufficientfor this purpose when proposing; there is a small (< 10%) decline in effective area foroff-axis pointing due to vignetting, but this is not sufficient to change the optimal strategy.
Observers should also be aware of the impact of uncertainty in the spectral shape on the
expected count rate and the uncertainty in the PSF wings to the Resolve illumination
pattern for the off-axis pointing. Note that values quoted here are based on the PSP
algorithm, measurements from Hitomi on the time needed to process events of each grade,measurements of the cross-talk contamination, and ground calibration measurements of
the PSF and the energy resolution. As more on-orbit XRISM measurements are made, themodel will be updated in future cycles.

Proposers’ Observatory Guide

79
Table 8.1: Resolve Count Rates

Flux
No filter
ND Filter
ND+Off-axis(0.3-10 keV 10−8 Total Hp+Mp Clean Total Hp+Mp Clean Total Hp+Mp Cleanerg s−1 cm−2 )(ph s−1 )(ph s−1 )(ph s−1 )
0.5
35
30
15
10
10
5
10
10
5
1.0
70
45
30
20
15
5
20
15
10
2.0
140
60
50
30
30
15
30
25
20
3.0
210
70
50
50
35
15
50
35
25
4.0
280
70
50
70
45
30
60
40
25
5.0
350
80
45
90
50
30
80
40
30
6.0
420
85
40
100
50
45
100
40
25
7.0
490
85
30
120
60
45
110
50
25
8.0
560
80
25
140
60
45
130
50
25
9.0
630
80
25
160
65
50
150
50
25
10.
700
80
10
170
70
55
160
50
30
20.
1410
70
1
340
80
40
330
65
30
50.
3520
30
0
860
80
5
810
65
35
All values assume an on-axis source except for the “off-axis” set at a fiducial offset of +0.5 arcmin
offset in both RA and Dec. “Clean” events are those H+Mp events whose average energy
resolution is within 2% of a fiducial 5 eV resolution. Bold text indicates the optimal rate andthus observing configuration. The input spectrum is a Γ = 2 power law with photoelectric
absorption from 1021 cm−2 of foreground atomic gas. Exact values are sensitive to the spectrum,
structure in the PSF wings, etc. This table is for illustration only.

Secondly, proposers should estimate the sensitivity of their scientific goals to smallchanges in energy resolution. We defined clean Hp+Mp events as those originating inpixels with average energy resolution within 2% of the native resolution (i.e., 5.1 eV fora native 5 eV). This would make the cross-talk degradation smaller than the calibration
uncertainty to the point where it could essentially be ignored. Proposals that are notfocused on fine line structures or that can accept greater velocity uncertainty may maximize
the total Hp+Mp rate rather than the “clean” rate. The clean rate cannot be estimatedfrom the ARFs alone but requires knowing the expected count rate in each pixel (e.g., fromheasim), so that one can make an estimate of the cross-talk rate in each pixel, and also theincident spectrum, so that one can estimate the energy degradation for a line of interest.
A tool to estimate the impact of cross-talk for an input spectrum will be made available
at https://heasarc.gsfc.nasa.gov/docs/xrism/proposals/index.html, but it is notnecessary: Table 8.1 can be used as a guide for proposing to maximize the total Hp+Mp
rate or the clean Hp+Mp rate.

8.5Filtering Cross-talk on the Ground

Once the data are in hand, the main objective is to maximize the clean Hp+Mp counts.
The definition of “clean” very much depends on the science goal: extracting a spectrum
from the entire array makes sense for maximizing the signal-to-noise ratio (S/N ) in a shortProposers’ Observatory Guide
Be filter
Neutral Density filter

200
Effective Area [cm2]80
default effective area

150

100

50

0

2

3

4
5
6
Energy [keV]

7

8 9 10Figure 8.3: Resolve effective area curves with the gate valve closed for no filter (black), the Be filter (red),and the neutral-density filter (blue). The neutral-density filter reduces the count rate by about a factor
of four across the bandpass. For filter transmission curves see Figure 5.13. The Be filter is optimized
for suppressing (continuum) events at low energy in favor of preserving high-resolution events in the Fe-K
complex.

time frame, while detailed velocity measurements using an Fe xxvi absorption line may
require selecting events only from pixels largely unaffected by cross-talk. We defined clean
above as 2% of the native energy resolution (i.e., 5.1 eV for a 5.0 eV native resolution),
which is well within the calibration uncertainty, but observers can estimate the cross-talk
impact from the data themselves and adopt their own definition of clean.
The simplest way to eliminate cross-talk contamination is to exclude ±25 ms from eachpulse in each electrically neighboring pixel, which can be captured via a good-time interval
file. This exclusion window is in the CALDB and, while not automatically executed in thepipeline processing, can be implemented with the XRISM software tools. The impact on
the effective exposure time is a strong function of pixel, since the pixels near the core of the
PSF will have by far the highest count rates, and their electrically neighboring pixels are
not only the ones physically adjacent on the detector. If this removes too many events, thenext simplest approach is to use the count rate per pixel to determine which electricallyneighboring pixels will be most affected by cross-talk degradation and omit those pixels
from a combined spectrum. This can be predicted as a function of count rate, position,
and a rough idea of the spectrum, as shown in Figures 8.2 and 8.4. For these purposes, the
MXS can be treated as part of the same source as it will be affected by cross-talk filtering.
We note that these predictions carry inherent uncertainty of about 1-2% of the native
energy resolution. This is because only a small fraction of cross-talk pulses significantly
affect the energy measurement (those that occur near the peak of a real pulse), and themagnitude of the error also depends on the relative energy of the real pulse and the cross-Proposers’ Observatory Guide

81

Figure 8.4: These six plots show a mock image (top left), raw incident count rate (top center), processed
Hp+Mp rate (top right), cross-talk child pulse rate (bottom left), effective spectral resolution (bottom
center), and centroid shift (bottom right) in each pixel for a 10−7 erg s−1 cm−2 source observed no filters.
The top right quadrant is sacrificed to maximize the Hp+Mp rate in the other three quadrants, as eachquadrant is processed separately in the PSP.

talk parent pulse. Since photon arrival is a Poisson process, any given observation will
deviate from the average a small amount.
To determine the importance of cross-talk contamination, consider that line centroids
can be measured to a precision
∆v
1
∆E
=
c
2.35(S/N ) E(8.1)

For an instrumental resolution of 5 eV, an Fe Kα line at 6.4 keV, and S/N = 10, one canmeasure ∆v to ±10 km s−1 . Degrading the resolution to 6 eV would lead to a constraintof ±12 km s−1 . Line broadening can be determined to a precision
2 + σ2σline∆σline
1
inst
=√
2
σline
σline
2(S/N )

(8.2)

For a line of intrinsic width 1 eV and an instrumental resolution of 5 eV, S/N ∼ 12 isneeded to measure σline to 30% precision. At 6 eV, S/N ∼ 15 is needed to make theProposers’ Observatory Guide82

same measurement. It is also important to consider the limiting precision on a prospective
measurement from systematic calibration uncertainty, such as gain differences between
pixels, which effectively degrades the energy resolution.

8.6

Summary of Observing Recommendations

The optimal combination of offset pointing and/or filter depends on the source type andscience goals. Here we provide basic guidelines for observing point sources of various
fluxes with the goal of characterizing Fe K emission or absorption above 6 keV in a source
otherwise characterized by an absorbed power law.
Table 8.2: Bright Source Observing Recommendations.
Flux
Recommendation
Clean Rate
(0.3-10 keV 10−8 erg s−1 cm−2 )(ph s−1 )0.5-1
On axis, no filter
20-30
1-3
On-axis, no filter
40-55
3-10
On-axis, ND
40-60
10-30
On-axis (0-0.5′ ), ND
40-50
50
Off-axis (1.0′ ), ND
30“Clean” events are those H+Mp events whose average energy resolution is within 2% of the fiducial 5 eV resolution. The input source is a point source with a a Γ = 2 power law spectrum withphotoelectric absorption from 1021 cm−2 of foreground atomic gas. Exact values are sensitive tothe spectrum, structure in the PSF wings, and off-axis position. At all fluxes, the appropriate
Xtend mode is 1/8 window+burst, as the source always produces >500 counts s−1 .

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