Chapter 4:

4.1.

X-ray Imaging (Spectroscopic) Telescope

X-ray Telescope for the Overall Solar-C/Plan-B Science

In addition to its own scientific standpoint, the X-ray Imaging (Spectroscopic)
Telescope shall have close scientific synergy with other Solar-C/Plan-B instruments tomaximize outcome from the Solar-C observatory as a whole. While SUVIT and EUVS
chiefly observe lower layers of solar atmosphere, the X-ray Imaging (Spectroscopic)
Telescope observes the uppermost layer of the Sun, the corona, to achievecomprehensive understanding on magneto-hydrodynamic processes across the entire
layers of solar atmosphere with various degrees of ionization. Science with
Solar-C/Plan-B relevant to the corona shall address (a) forms and mechanisms ofstorage and dissipation of energy that was transferred upwards through photosphere andchromosphere, (b) quantitative understanding on the physics of magnetic reconnection
which is believed to be the central engine for converting magnetic energy to thermal
and kinetic energy of coronal plasmas, and possibly, (c) diagnostics on wave
phenomena along magnetic field lines. Investigation of mechanism(s) for explosiveenergy release may also have relevance to space weather studies.
In pursuing such science capabilities, we consider the following two possibilities for
the Solar-C/Plan-B X-ray Imaging Telescope:(1) Photon-counting imaging spectroscopy X-ray telescope(2) Ultra-high spatial resolution normal incidence EUV telescope
The photon-counting telescope adopts grazing-incidence optics and aims to perform
imaging-spectroscopy of the soft X-ray corona from below 1 keV up to ~5 keV with its
main scientific target being quantitative investigation of dissipation processes that act in
the corona such as heating of active region corona and magnetic reconnection in flares.
The upper-end of the energy range can be increased up to ~10 keV if we adopt a smaller
grazing-incidence angle (~5 keV with 0.9 degrees while ~10 keV with 0.45 degrees) as
will be described in the subsequent part of this report. The energy coverage up to ~10
keV will be particularly powerful in investigating heating and acceleration of plasmas in
the reconnecting magnetic field structure in flares.
While the photon-counting telescope should be powerful in studying energy
dissipation processes in the corona, it may have less overlap with the other
Solar-C/Plan-B instruments (SUVIT and EUVS) in terms of temperature coverage(hence less overlap in the atmospheric layers of the Sun) and spatial resolution becauseof its grazing-incidence optics. Taking this point into account, the other possibility, theultra-high spatial resolution normal incidence telescope has been considered. The
normal incidence telescope puts its emphasis on delivering complementary observables
for EUVS, providing context imagery information for EUVS with higher spatial

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resolution than EUVS, with a wide field-of-view and with overlapping temperature
coverage thus aims to investigate connection between transition region and the low
corona.

4.2.
4.2.1.

Photon-Counting Imaging Spectroscopy X-ray Telescope
Scientific Objectives

The photon-counting imaging spectroscopy X-ray telescope will provide thefirst-ever opportunity to perform imaging-spectroscopic investigation of the corona insoft X-ray wavelength range, from below 1 keV up to ~5 keV (Figure 4.1).

Figure 4.1. Status of past imaging-spectroscopic investigation in EUV/X-ray wavelengths overlaid withmodel solar spectra synthesized with CHIANTI Ver. 6. Imaging-spectroscopic investigation of the soft X-ray
corona has so far remained totally unexplored.

Among Plan-B science goals detailed in Section 2, two main scientific targets of thistelescope with photon-counting capability is summarized as follows:(1) Quantitative investigation of structure and dynamics of magnetized plasmas
around the site of magnetic reconnection.
It is widely perceived, based on imaging observations with Yohkoh, that magnetic
reconnection is the engine for liberating magnetic energy for plasma heating andparticle acceleration in solar flares. Nevertheless, one has to say that understanding of
the reconnection process, such as structures of slow and fast shocks, and temperaturesand heating status of upstream and downstream plasma flows around the shocks, still

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remains almost unclear. The most critical piece of information that has been missing totackle this problem is spatially-resolved spectral information in the reconnecting
magnetic structure. Previous X-ray telescopes (Yohkoh SXT and Hinode XRT) could
only provide a single 'representative' temperature in the form of filter-ratio temperature
which is more like a mean temperature weighted by the differential emission measure(DEM) distribution over the line-of-sight plasmas (Acton, Weston, and Bruner 1999).
This feature of filter-ratio temperatures tends to smear out any spectral features acrosscoronal structures unless the bulk of plasmas (namely, those with the largest fraction of
DEM) changes its spectral shape. The photon-counting imaging-spectroscopy ofreconnecting magnetic structure will characterize plasma status upstream anddownstream of shock structures associated with magnetic reconnection. This would, inturn, identify location and structure of the reconnection shocks. Spectral information
with low-scattering focusing optics for a variety of flares should provide key clues inquantitatively understanding the physics of magnetic reconnection.
While the photon-counting X-ray telescope can derive electron temperatures as thecontinuum Bremsstrahlung emission, emission-line diagnostic with EUVS can deliver
information on ion temperatures. One remarkable possibility is that the combination of
spatially-resolved electron and ion temperatures will provide us with significantly
powerful tool for studying electron and ion heating especially in the downstream of theshocks.
(2) Investigate temperature structure of active region corona, in particular that forhigh temperatures (above ~5 MK)
One of the key information necessary for unravelling the mechanism(s) of coronalheating, in particular the hot corona around active regions, is to identify evolution ofspectral structures together with their distribution across lateral (across active regions)or in vertical (hight distribution) extents. This will enable one to locate site of heat input(hence dissipation site) across the area and subsequent heat flow that results ingenerating the hot corona. This approach may be particularly effective if the heating isintermittent (i.e., by nano-flares). The non-dispersive nature of photon-counting
imaging-spectroscopy, in contrast to dispersive, scanning, instruments will be effectivein identifying such heat input and subsequent flow in the corona by measuring spectral
evolution taken simultaneously across a wide field-of-view. Instruments that utilize
EUV wavelengths may not be suitable for this investigation due to lack of strong EUV
emission lines at high temperatures from active regions.
Discovery space with imaging spectroscopy of an extended energy range
In this report, the X-ray grazing-incidence angle to the mirror of 0.9 degrees set as thebaseline. While this grazing-incidence angle provides certain capability for performing
imaging-spectroscopic observations of active regions (see below for details), the highest
energy available for flare observations would be limited to ~5 keV (Figure 4.13). An
interesting possibility which greatly enhances observation capability for flares, withdecreased capability for active regions, is to go to an even smaller incidence angle. For
example, the grazing-incidence angle of 0.45 degrees will increase the highest energy

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available for flare observations to as high as ~10 keV (Figure 4.14).
The small incidence angle (e.g., 0.45 degrees) for the photon-counting telescope will
enable imaging spectroscopy in the energy range of a few to ~10 keV. This performanceprovides the following discovery space for flare sciences in the field of reconnection
study:(1a) Identification of high-temperature plasmas in the flaring magnetic field structure
The increase in the energy range up to ~10 keV will enable one to observe Fe (~6.7
keV) and Fe/Ni (~8 keV) line complexes together with thermal-bremsstrahlung
continuum from high temperature plasmas in energy spectra of flares. When combinedwith the expected imaging-spectroscopic performance of the telescope shown in Figure
4.15, this means that we'll be able to track spatial distribution and its evolution ofsuper-hot (> ~30 MK) plasmas with angular scale and temporal resolution combinationof 2" with every 5 s (or 1" with every 20 s). Imaging spectroscopy with such fine spatial
scale and temporal resolution will enable us to identify site of in-situ heating and to
discuss heating of plasmas in the reconnecting magnetic field structure.
(1b) Evolution of non-thermal (or supra-thermal) electron population
The extended energy coverage should also provide breakthrough in understanding
particle acceleration mechanism(s) in solar flares. It is generally believed that initial
particle acceleration takes place inside the reconnection region (forming reconnection
outflows), followed by transport to and further acceleration somewhere at the outer partof it (see Tsuneta & Naito (1998) for the case of fast-shock acceleration). However, it
has so far remained unclear, or even unexplored, where and how particles (ions andelectrons) are accelerated in the reconnecting structure. While soft X-ray imagingspectroscopy of the reconnecting magnetic structure should naturally be powerful in
identifying, and tracing the evolution of, e.g., shock structures associated with magnetic
reconnection, observations with the extended energy range of up to ~10 keV (ascompared to ~5 keV in the baseline design) would provide enourmous leap inunravelling the nature of particle acceleration in flares. This is because the observations
will deliver information on how and where supra-thermal electron distributions(manifest themselves as power-law X-ray spectra) grow out of the ambient thermal
electron distribution, especially near the onset of particle acceleration where soft X-ray
emission from thermal plasmas along the line of sight is yet to be dominated. As themean energy of thermal electrons is typically 2-3 keV, X-ray emission by supra-thermal
electrons can be well identified in the ~5-10 keV range. Note there has been a numberof observations reported since 1980's that non-thermal X-ray emission from flares may
extend down below 10 keV. Furthermore, there is related observational support from insitu observations of the Earth's magnetosphere that supra-thermal electrons appear inenergy spectra around 4 times as large as thermal energy (2 keV) near the reconnection
region (see, e.g., Imada et al. 2007). Thus, the extended energy coverage, coupled withhigh dynamic-range imagery with soft X-ray focusing optics, will give us a unique
opportunity to locate the site of electron acceleration and to study the nature of theacceleration in relationship with the reconnecting magnetic field structures such as

4-4shocks.
Table 4.1. Key parameters of the photon-counting X-ray telescope.
Item

Description

Overall description

Perform imaging investigation of the X-ray corona withthe following two modes. Two modes selectable (notsimultaneously).
(1) Photon-counting mode: Imaging spectroscopy for alimited region of interest.
(2) Photon-integration mode: Imagery with the entire,or part of, the imaging array.
Walter Type I (-like) grazing-incidence telescope with asegment mirror. A CMOS-APS detector to be used as the
focal-plane imaging array.

Envelope dimensionsof the telescope~40cm × 40cm × 4.5m (TBD)

Angular resolution&
Temporal resolution
for photon counting

For photon-integration mode:~1" (~0.5"/pixel)
For photon-counting mode:
For active regions and flares, generate an X-ray
spectrum for spatial area whose size*~< 2"×2" (goal) / ~4"×4" (min.)within every< 10 s (goal) / 30 s (min.)For photon-integration mode*:~13'×13'
For photon-counting mode**:>~3'×3' (goal)~80"×80" (nom.)~40"×40" (min.)'Elongated' FOV (such as ~40"×13') is possibile

Field of view

Energy range

Remarks~0.5-5 keV* (baseline)~0.5-~10 keV* (under study)

Energy resolution

Data rate∆E (FWHM)*:~130 eV (goal), 150 eV (nom.) @ 5.9 keV~74 eV (goal) @ 1.5 keV
Readout noise:~43 eV FWHM (goal)(With the event processing in Section 6.2.3.4)
max. 110 Mbps for 512×512 photon-counting area
max. 6.9 Mbps for 128×128 photon-counting area

Telescope plus the
focal-plane camera.
Includes housings.
* Spatial area size inwhich a spectrum withgood counting statistics
can be synthesized withreasonable integrationtime.

* Off-axis image blur nottaken into account.
** Area with whichphoton-counting analysis
can be performed.

* Energy range above
2 keV available during
flares, with photoncounting mode.
* Includes readout noise.

No data compression
assumed. For 10-bits A/D
conversion.
Entrance filter for rejecting
visible light and heat input
to the telescope.

Pre-filter

Same as Hinode/XRT.

Visible-light optics

TBD

X-ray flux monitor

TBD; for controlling X-ray flux attenuation filters.

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Besides photon-counting capability, the telescope will also be capable of providing
images of soft X-ray corona (with conventional photon-integration, or 'flux mode') withthe highest angular resolution of 1 arcsec (0.5 arcsec pixel size) as a solar X-ray imager.
Such images will deliver context information for other Plan-B instruments.
By adopting an X-ray CMOS active-pixel sensor (APS) device for the focal-plane
imaging array, the photon-counting X-ray telescope will be free from "spill-over" ofsignal charges for over-exposed images which is often seen with soft X-ray imagers
employing CCDs without anti-blooming pixel architecture, including Hinode XRT and
Yohkoh SXT. The spill-over-free feature of the telescope, combined with highdynamic-range focusing optics imagery, should be particularly powerful for the soft
X-ray corona, in detecting faint features that could be present right next to a brightfeature, such as material flows and ejecta associated with magnetic reconnection.
With the advent of recent progress in high-speed, low noise, CMOS-APS detectors,
the photon-counting X-ray telescope for Solar-C/Plan-B should provide breakthrough in
our understanding on energy storage and dissipation in the corona, including the physics
of magnetic reconnection. Key parameters of the photon-counting telescope are given in
Table 4.1.

4.2.2
4.2.2.1

Required Specification for the Instrument
OpticsPhoton-counting imagery in X-rays needs be done with grazing-incidence optics thatcan cover, continuously, a wide spectral range of X-ray emission from the corona.
Normal-incidence imagers may not be suitable for the photon-counting purpose because(1) they can image only a narrow wavelength range for each passband, and (2) since thewavelengths for normal-incidence imagers are in the EUV range whose photon energies
below ~100 eV, silicon detectors such as CMOS-APS, whose best-expected FWHM
energy resolution would be around ~50 eV at those energies, are not sufficient forresolving fine spectral features in the EUV range. A grazing incidence telescope ischosen to utilize the sensitivity to the hotter plasmas and the associated spectral lines, insoft X-rays.
4.2.2.2

Angular Resolution

It has been suggested by recent Hinode EIS observations (Hara et al. 2008) that
sub-arcsecond structures in coronal loops may be important in understanding
elementary processes of coronal heating for active regions. In this regard, pixel size ofless than 1", say, 0.5", would be preferable for photon-integration imagery. For
photon-counting imaging spectroscopy, while angular resolution of, e.g., 5" (that is,energy spectra with every 5" square; half-resolution pixel size of Yohkoh SXT) could
still provide new spectroscopic information on the hot corona, especially for large
magnetic structures such as reconnecting loop structure seen in long-duration flares(LDEs), higher angular resolution even beyond that (e.g., 2" square or even higher)
would be much desired in investigating spectral features in active region loops and in

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reconnecting magnetic structures in flares.
4.2.2.3 Temporal Resolution / Cadence
From photon counting viewpoint, capability of generating an energy spectrum every
30 s or shorter should be necessary for active regions to see their thermal evolution. For
flares, temporal resolution for generating energy spectra of 10 s, or shorter, should bedesired.
4.2.2.4

Energy Range and Energy Resolution

Imaging-spectroscopic capability for hot corona should be essential for investigating
flare science such as reconnecting structures. This leads to a requirement that thetelescope should be sensitive beyond 2 keV (which is the upper limit for Hinode XRT),
to several keV's. Also, sensitivity below 1 keV is desired for imagery with photon
integration to image quiet Sun and coronal holes. Regarding energy resolution forphoton-counting observations, FWHM resolution of ~150 eV is needed to observe
higher end of the energy range (several keV's). For low-energy range of around 1 keV,necessary energy resolution is currently under investigation. As is discussed in Section
4.2.4, a combination of 5 electrons (rms) read-out noise and Fano-limited energy
resolution (with Fano factor F = 0.12 for silicon) provides comfortable energy
resolution for observing emission lines in the ~1 to 2 keV range.

4.2.3
4.2.3.1

Baseline Instrumentation
Overview

In response to specifications based on scientific requirements, we propose aphoton-counting soft X-ray telescope whose observational performance summarized inTable 4.1. In this Section, we explain key items for the proposed telescope.
(a) Optics System
Figure 4.2 indicates a preliminary sketch of the grazing-incidence telescope withphoton-counting capability. The telescope consists of (1) a pre-filter in front of theentrance aperture, (2) front-end grazing-incidence mirror of Walter Type I (-like) design,(3) focal-plane CMOS image sensor, (4) analysis/attenuation filters in front of the
CMOS detector, (5) focal-plane shutter, and (6) camera control & event processing
electronics (items (1), (5) and (6) not shown in Figure 4.2). There may be an aperture
adjustment mechanism under consideration that is to be mounted in front of thetelescope aperture to reduce incident flux to an appropriate level when performing
photon-counting observations. The telescope will have ~4 m of focal length.

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Figure 4.2.

Schematic illustration of a preliminary photon-counting telescope concept.

(b) Event Processing System
In addition to the optical performance of the grazing-incidence mirror, accurate
measurement of energy and position on the detector for each detected photon, withavoiding significant pulse pile-up effect, is a crucial key for the success of this
instrument. Figure 4.3 presents a preliminary diagram for event processing underphoton-counting measurement.
Besides the function blocks indicated in Figure 4.3, the telescope should employ
appropriate attenuation filter or aperture size in response to the incident flux level.

Figure 4.3. A preliminary block diagram describing the event processing scheme for detected X-ray
photons under photon-counting measurement.

4.2.3.2

Optics/Mirror

A piece of grazing-incidence segment mirror, consisting of 1/8 of an entire circle, will
be adopted for the X-ray mirror of the telescope. The mirror employs a Walter-I (-like)design with 2 reflections, with the length (in optics axis direction) of the first reflection
segment designed to be 40 mm (TBD). While the radius of the mirror r (radius of an
entire circle along which the segment mirror is laid), located at f = 4 m in front of thedetector, changes with different grazing incidence angle to the mirror θi with thefollowing relationr = f × 4θ i ,4-8

we take θi = 0.9 degrees as the baseline to perform imaging-spectroscopy both for active
regions and flares. Grazing angles of 0.45, 1.35, and 1.8 degrees have also been studied.
Besides the baseline incidence angle, θi = 0.45 degrees could be an interesting option
for flare sciences (Section 4.2.4.3) but with penalty of poor observational performancefor active regions. The baseline diameter of the mirror with the grazing incidence angleof 0.9 degrees is 50.3 cm.
Unlike Hinode XRT, the mirror must be coated with iridium (Ir) to attain highreflectivity, hence sensitivity of the telescope, for keV X-rays.
A conceptual schematics of the segment mirror mount is shown in Figure 4.4, with adrawing for the mirror-detector package in Figure 4.5.

Figure 4.4. Left: Close-up of a segment mirror unit consisting of primary and secondaryreflecting surfaces, having 1/8 segment of the entire circle. Right: Preliminary conceptual
illustration on the focus mechanism which moves the mirror segment along the direction of the
optics axis.

Figure 4.5. An illustration indicating segment mirror unit (left) and the focal-plane detector (right).
Ray paths for X-rays from the Sun also shown in the figure.

A preliminary set of baseline parameters for the X-ray mirror is summarized in Table
4.2. A key requirement to the grazing-incidence mirror is that it should keep very small
figure error (deviation from ideal mirror surface profile) to achieve 1 arcsec resolution(0.5 arcsec pixel size) while maintaining small micro-roughness for suppressing X-ray

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scattering from other (bright) features on the Sun. A first-cut, rough estimate indicatesthat the figure error of ~10 nm P-V (TBD) would be needed to have Half-Power
Diameter of 0.5 arcsec for angular resolution and micro-roughness of 3-4 Å r.m.s.
(TBD) to suppress off-axis scattering wing level down to ~10-5 of the core at 1 arcmin
off-axis at Al-K.
While fabrication of a mirror with the above specification is challenging,development of the mirror with Japanese domestic partners is currently underinvestigation which might provide possible solution for the fabrication.
Table 4.2.

Baseline parameters for the X-ray mirror (preliminary).

Item

Description

Remarks

Mirror type

Walter Type I (-like) grazing incidence mirror, with
two reflections at the primary and the secondaryreflection segments.
40 mm for the primary reflection segment.
TBD mm (~50 mm) for the secondary reflection
segment.

TBD; May be increased
depending on tolerance
for abberation.

Length of thereflection segment
Annular segment
profileGrazing incidence
angle
Annular radius ofcurvature

1/8 of the entire circle.

Mirror surface

Coated with iridium (Ir)

Focal length

4.0 m

Plate scale

0.41"/8µm (0.5"/10µm)

0.9 degrees

Baseline value.

251.3 mm

For grazing-incidence
angle of 0.9 degrees.
Coating is a must forscience performance ofthe telescope.
8-10 µm corresponds tothe pixel size of thebaseline detector.

Angular resolution
performance
Core
Wing

4.2.3.3

Figure error of the mirror ~10 nm P-V (TBD) to have
0.5 arcsec HPD (Half-Power Diameter).
Mirror micro-roughness 3-4 Å r.m.s. (TBD) tosufficiently suppress off-axis scattering wing level
down to ~10-5 of the core at 1-arcmin off-axis at Al-K.

Focal-Plane Detector

Table 4.3 summarizes baseline specifications for the focal-plane detector to meet the
instrument performance tabulated in Table 4.1. In order to achieve high frame readout at
the focal-plane detector, we identify CMOS image sensors as the most appropriate forour purpose. There are several active-pixel detectors under development besides CMOS
which aim pixel-based spectroscopy on X-ray photons, such as DEPFET for IXO X-ray
observatory (former XEUS; Strüder et al. 2003, 2004, Treis et al. 2005, Strüder 2008)
and pixel detectors for high-energy particle accelerator applications such as for LHC

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(e.g., Llopart et al. 2007), these detectors mostly have large pixel size of typically ~50µm, with the smallest one at most ~25 µm. The large pixel size of these detectors,
unfortunately, is not consistent with the required angular size of a pixel of ~0.5" unlessthe telescope has significantly large focal length exceeding 10 m.
Recently, there has been remarkable development in CMOS active-pixel image
sensors which appear to be best suited as our photon-counting detector because of theirsmall pixel size (can easily be smaller than 10 µm) and their potential of high frame
readout capability.
Table 4.3.

Baseline specification for the focal-plane detector (preliminary).

Item

Description

Remarks

Device type

CMOS-APS

Device from e2v considered
to be the baseline.

Format2k×2k or larger

Pixel size

8-10 µm

Photon-counting area

512×512 (goal)
100×100 (minimum)

Frame readout rate
for photon counting

1000 fps

Pixel gain
Low-energy X-ray
sensitivity
Readout noise
Energy resolution

Should be changeable overthe imaging area.

>a few × 10 (>50) µV/e- for photon-counting pixels
a few µV/e- for photon-integration pixels
Back-illuminated device, or thin (~<1 µm) SiO2 absorption
layer if a front-illuminated device is employed.
5 e- rms (TBD; goal)≤50 eV FWHM (minimum)
Fano-limited (F=0.12) (goal)≤150 eV FWHM (minimum)

Charge spreading

At least 3-σ electron cloud diameter smaller than 3 pixels.

Pulse-height A/Dconversion≥10 bits

Should be switchable overthe imaging area.

Figure 4.6. Envisaged CMOS-APS for WFS whose development ongoing at e2v. An area of
1680×1680 pixels is planned to be read out with a frame rate of 700 fps (nominal) and with readout
noise of ~3 e- r.m.s. Figure courtesy of e2v.

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As of this writing, the most promising series of devices along the line of our need isthose which could be provided by e2v Technologies, U.K. A large-format wavefront
sensor (WFS) for night astronomy is currently under development by e2v (Figure 4.6).
The target specification for the WFS is to read out 1.6k × 1.6k pixels (24 µm pixel size)area at a frame rate of 700 fps (nominal) with low readout noise (a few e- r.m.s.),backside illumination with bulk depletion using high-resistivity epitaxial silicon. They
have recently successfully completed demonstration of their approach on pixel
architecture to achieve low readout noise together with high frame rate, and are now
quite confident that most, if not all (except the largest case of the photon-counting area),of baseline specifications in Table 4.3 can be met with their pixel architecture developed
for the WFS. Figure 4.7 depicts possible readout of the photon-counting area among the
2k × 2k imaging area (pixel size 8-10 µm) based on the existing architecture at e2v withthe background of WFS heritage. Final pixel size will be determined considering effectof charge spreading by X-ray photons on event detection efficiency.
The extended readout area in the horizontal direction in the figure would beparticularly effective when considering the possibility of flexible spacecraft roll to bementioned in Chapter 5 in such a way that the horizontally-extended area can serve as a
vertual 'slit' along which imaging spectroscopy can be performed.

Figure 4.7. Possible readout scheme of photon-counting area in the 2k×2k array with the existing
pixel architecture at e2v. In the figure, areas labelled as '(u)' are read out from the upper side of thedetector while those with '(d)' from the lower side, each with 1000 fps. [A] 2k×200 area in the middleof the imaing area to be read out with 1000 fps. [B] Two separate 2k×100 areas to be read out. [C]
Switch read out of area (d1) and (d2) each other so that a total of 2k×200 area in the lower half of thedetector be read out at 500 fps, which may be useful for observing weak sources extending large
spatial extent. Figure courtesy of e2v.

In parallel to the development of high frame-rate readout, back-thinning of
CMOS-APS is also being extensively promoted at e2v, using high resistivity (~1kΩ-cm) epitaxial silicon to have the bulk silicon depleted instead of depleting only alimited portion around the p-n junction of the photo-diode in each pixel thus attaining adepletion layer thickness of ~10 µm or even more. A sample of such back-thinned
CMOS-APS is said to become commercially available next year (2011). A prototype
WFS device with ~800×600 pixels and with a readout rate of 700 fps (1000 fps max.)
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for the entire area is scheduled to appear in 2012, again with back-thinning and withhigh-resistivity silicon.
As for spreading of signal charges by incident X-ray photons, only limited results
have so far been available with e2v CMOS-APS devices. Nonetheless, results from
X-ray illumination with Mn target (that emits Mn Kα and Kβ lines besides
Bremsstrahlung continuum) for a front-illuminated device performed at Centre forElectronic Imaging at The Open University indicates the X-ray signals are contained
within 3×3 pixels area with ~20% of Mn Kα (5.9 keV) events resulting in single-pixel
events even with a small pixel size of 5.8 µm and with low resistivity silicon (hence
shallow depletion layer). This result suggests promising prospects that the extent ofcharge spreading (and also fraction of single-pixel events) would even be better with alarger pixel size of 8-10 µm and with high resistivity silicon.
To summarize, at this very moment there is no CMOS-APS device present that
immediately meets the baseline specifications for the photon counting. However, withthe rapid development of CMOS-APS currently ongoing towards the desired directionof (1) high-speed and low-noise readout capability and (2) good soft X-ray
spectroscopic capability, the photon-counting detector is expected to become even more
realistic within the upcoming 1-2 years.
In the subsequent section of this Report, we assume back-illuminated CMOS device
from e2v, with 10 µm-thick depletion layer and with 8 µm pixel size for assessing
science capability of the photon-counting telescope.
4.2.3.4

Flux Attenuation System

For performing photon-counting observations for flares and active regions, it isdeemed necessary to employ attenuation system for the incident X-ray flux so as to
avoid event pile-up on the focal-plane detector. Either a set of attenuation filters
mounted on a filter wheel(s) or an aperture adustment mechanism would be employed
for the flux attenuation system with the former may be more suited to observe higher
end of the spectrum because the attenuation filters absorpt lower-energy X-ray photons.
While logic for selecting certain attenuation filter or a particular aperture size (withprobably using signals from a separate flux monitor) needs be elaborated, observations
with the baseline geo-stationary orbit (see Section 5) could provide a remarkable chanceof performing such attenuation on true real-time basis.

4.2.4
4.2.4.1

Expected Capability
Introduction

Figure 4.8 shows effective area profiles of the telescope for the four
grazing-incidence angles θi studied (θi = 0.45, 0.9, 1.35, and 1.8 degrees), with θi = 0.9
degrees being considered as the baseline. The effective area profiles include (1) X-ray
transmission of the pre-filter (same material design as Hinode XRT, namely, 1200 Å of

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aluminum on 2500 Å polyimide assumed), (2) reflectivity of the grazing-incidence
mirror, (3) projected QE of the focal-plane CMOS detector, and (4) geometrical area ofthe telescope that is determined by the aperture size of the first reflecting segment of themirror, when viewed from the Sun.

Figure 4.8. Effective area profiles for grazing incidence angles of 1.8 (black), 1.35 (red), 0.9 (green),
and 0.45 (blue) degrees. The effective area profile for Hinode XRT (without Ir coating on the mirror
surface) is also shown for comparison (black dashed line). QE profile for the back-illuminated e2v
CMOS device assumed for all the four grazing incidence angle profiles.

Figure 4.9. Left: X-ray count spectra for a model active region with CHIANTI V6.0, for four
grazing-incidence angles, 1.8, 1.35, 0.9, and 0.45 degrees (in black, red, gree, and blue, respectively).
No attenuation filters assumed to be present. Right: Those for a model flare (GOES class M2). Again
without attenuation filters.

Figure 4.9 indicates X-ray count spectra for a model active region (left panel) and amodel flare (right panel) from CHIANTI V6.0 (Dere et al. 2009), without focal-planefilters. Readout noise and energy resolution indicated as "goal" in Table 6.2-4 are used
for generating these plots. Note with the back-illuminated CMOS neither active region
nor flare count spectra can be observed as they are in the photon-counting mode
because incident X-ray count rates will be too high (e.g., in the flare case, ~3×105
cts/s/pixel for θi = 0.9 deg.) to perform any meaningful photon counting avoiding pulse
4-14pile-up. For this reason, focal-plane attenuation filters should be necessary. The
attenuation filters will make count spectra to be actually measured different from thosein Figure 4.9, with low-keV portion to be significantly attenuated. In thephoton-integration mode, such attenuation filters will not be necessary but analysis
filters for performing filter-ratio analysis may be employed.
Signal count rates to the focal-plane detector of the telescope without attenuation
filters (nor focal-plane analysis filters) is summarized in Table 4.4. Here, the baseline
back-illuminated CMOS detector is assumed for estimating signal count rates.
Table 4.4.

Estimated signal count rate for the telescope, with the baseline detector and thegrazing-incidence angle of 0.9o.

Target(1)Flare (M2)

Active Region

Quiet Sun

Coronal HolePhoton count rate (cts/s/pxl)

3.1×105 (2)2.2×102 (2)

7.9

1.7

Electron rate (e-/s/pxl)

7.4×107

1.8×104

1.9×102

3.5×101

Exposure time for photon1.0 ms
4.2 s
3.9×102 s
2.1×103 s
integration(3)(1) Emission spectrum from each target created from CHIANTI V6.0.
(2) Attenuation filters (and/or aperture adjuster) need be inserted in order to perform photon-counting
spectroscopy.
(3) Time for generating 75 ke-, which is half the nominal full-well capacity of the CCD for
Hinode XRT.

4.2.4.2

Active Regions

Figure 4.10 presents simulated photon count spectra to be detected with the photoncounting mode of the telescope, with the baseline grazing-incidence angle of θi = 0.9
degrees. A model DEM for active region (Vernazza and Reeves 1978) with CHIANTI
V6.0 emission code is used for obtaining solar emission from a "typical" active region.
The spectra assumes 1024 energy bins between 30 eV to 10 keV, with binning made insuch a way that count rate from each summed-bin has at least 5-σ statistical significanceagainst Poisson errors (only photon counts are considered here as the source of theerror), with an imposed limitation that the maximum number of binning be 20. A set of
readout noise of 5 electrons r.m.s. and Fano-limited energy resolution (with Fano factor
F = 0.12) is adopted here to produce the spectra. This set of readout noise and energy
resolution is listed as goal for the detector specification in Table 4.3.
Left panel of Figure 4.10 shows a spectrum with 262 s net integration time while theright panel 654 s net integration. Here, net integration time stands for the total of
integration time from each pixel that contributes for generating a single energy
spectrum. For example, for generating the left-side spectrum of the figure, it will take16 s with 4×4 pixels (1.6" square).
As seen in the figure, an emission line complex of Mg may be detected around 1.34
keV even with the 262 s net integration case with positive, or marginal, indication ofanother emission line complex of Si around 1.85 keV. As the contribution curves forthese emission line complexes peak at 5-8 MK, any deviation of measured emision line
complexes from Figure 4.10 immediately give us information on distribution and

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evolution of high-temperature (≥5 MK) plasmas that has never been provided before.

Figure 4.10. Simulated photon count spectra for a "typical" active region with the baseline
grazing-incidence angle of 0.9 degrees and with the baseline CMOS detector. Left: With net integration
time of 262 s. Right: Net integration time 654 s.

If we concentrate on the low-temperature emission in the energy range below ~1 keV,
energy spectra in that energy range can be generated even quickly. Figure 4.11 shows anexample of active region spectrum with 2 µm-thick Be filter employed as theattenuation filter (yielding incident photon count rate of 39.1 counts/s/pixel to thedetector) with net integration time of 130 s. This means the spectrum in the figure
would be obtained every 14 s with 3×3 pixels (1.2" square).

Figure 4.11. Simulated photon count spectra for a "typical" active region with the baseline
grazing-incidence angle of 0.9 degrees and with the baseline CMOS detector, with net integration
time of 130 s.

Apart from thermal emission so far studied, if non-thermal emission were to extend
down to, say a few keV range, at the early phase of a flare, such non-thermal emission
is likely to be well detected as continuum component in the spectra, in a spatiallyrevolved way. Note detection of such non-thermal components would not be available
with EUV spectrographs since they manifest themselves as continuum emission, andcan be done only with the X-ray imaging-spectroscopic approach. So far, spatiallyresolved observations of non-thermal components down to a few-keV range have notbeen performed except for very limited cases (Krucker et al. 2002, Hannah et al. 2008).
4-16

Observations with the photon-counting X-ray telescope will provide crucial information
on physical condition of the corona that triggers explosive energy release in flares.
4.2.4.3

FlaresIn the case of flares (with a model flare of GOES class M2 studied with CHIANTI
spectral code), attenuation filters are necessary to reduce incident rate of X-ray photons
to the detector in order to avoid excess photon pile-up, down to a level of several times
10 photons/s/pixel. Energy spectra thus observed will therefore undergo significant
attenuation of low-keV energy components by the attenuation filters. Figure 4.12 shows
expected photon count spectrum profiles for a model flare of GOES M2 class to be
detected with the photon-counting mode of the telescope, for different grazingincidence angles, with an attenuation filter (beryllium assumed) inserted in each case to
reduce the incident photon count rate down to ~40 photons/s/pixel. As seen in the figure,peak position of the spectrum changes with different grazing-incidence angles. Note in
actual observations (with limited integration time for a spectrum) count rates down toroughly 2 orders of magnitude lower than the peak would give a practical lower limit
for meaningful spectra with good, or reasonable, photon statistics. Hence energy rangethat can be explored differs with different grazing-incident angles.

Figure 4.12. Photon count spectrum profiles expected from a model GOES M2 flare for the four
different grazing-incident angles. An attenuation filter made of beryllium assumed for each profile;
thickness 0.42 mm for θi = 1.8o, 0.8 mm for θi = 1.35o, 1.5 mm for θi = 0.9o, and 1.5 mm for θi =0.45o so that each spectrum has a similar total count rate of ~40 counts/s/pixel. Note the profiles are
ideal ones with infinitely-good photon statistics.

Figure 4.13 presents simulated photon count spectra with the photon-counting modeof the telescope, with the baseline grazing-incidence angle of θi = 0.9 degrees. A
focal-plane attenuation filter of 1.5 mm-thick beryllium is inserted to reduce incident
count rate down to 42.3 counts/s/pixel. A model DEM for a GOES M2-class flare (Dereand Cook, 1979) with CHIANTI V6.0 emission code is used for obtaining solar

4-17

emission for this flare case. For displaying flare count spectra, binning was made so thatcount rate from each summed-bin has at least 10-σ statistical significance against
Poisson errors (only photon counts are considered here as the source of the error), withthe maximum number of binning set to be 20. The left panel of the figure corresponds to
140 s net integration while the right panel 840 s. Even with the 140 s net integration
time (that means 15.6 s integration for a 1.2"×1.2" area (3×3 pixels)), signature of Ca
XIX emission (emission from 20 MK plasma) at ~3.9 keV can be seen in addition to thecontinuum thermal emission. Note these spectra only reflect thermal emission fromflare plasmas; any non-thermal emission, if present, can be overlapped to these spectra.

Figure 4.13. Simulated photon count spectra for a model flare of GOES M2 class with thephoton-counting mode of the telescope of θi = 0.9o. Left: Net integration time 140 s. Right: Net integration
time 840 s.

As supplementary information, Figure 4.14 shows simulated flare spectra withgrazing-incidence angle of θi = 0.45 degrees. This grazing angle could provide aninteresting possibility that, in addition to the Ca XIX emission Fe and Fe/Ni emission
complexes around 6.7 keV and 8 keV, respectively, may also be detected.

4.2.4.4

Imaging-Spectroscopic Performance

Figure 4.14. Simulated photon count spectra for a model flare of GOES M2 class with thephoton-counting mode of the telescope of θi = 0.45o. Left: Net integration time 133 s. Right: Net integration
time 800 s. 1.5 mm-thick Be filter assumed as the attenuation filter so that the incident photon count rate to
the detector reduced to 40.4 counts/s/pixel.

4-18Based on the above observation capability for active regions and flares,
imaging-spectroscopic performance of the photon-counting telescope is schematicallyillustrated in Figure 4.15. The figure presents angular size (in arcsec) of observation
target versus temporal resolution that can be attained with the baseline photon-counting
telescope. (Here, 'baseline' refers to instrument capability summarized in Table 4.2 and4.4.) Lower-left direction in the figure corresponds to smaller target (angular resolution)with higher temporal cadence hence preferable for observations. For performing

Figure 4.15. Imaging-spectroscopic capability expected for the baseline photon-counting telescope. See
text for details.

imaging-spectroscopy on high temperature emission in the range of 1-2 keV, upper-rightside region from the solid blue line (with label 'Active Region (high-T)') can beaddressed by the instrument, while angular size-temporal resolution combination in theleft-side region would result in insufficient photon statistics for generating soft X-ray
energy spectra. On the other hand, for imaging spectroscopy on low temperatureemission in the range below 1 keV, upper-right side region from the blue dashed line(with label 'Active Region (low-T)') can be addressed.
For flare observations, the border is shown in solid red line (with label 'Flare (M2)'),and imaging-spectroscopic observations can be performed with good photon statistics
for the region upper-right of the border. If we could have small charge spreading (lessthan 3×3 pixels), or could adopt a better event selection method (if it exists) other thanthe one discussed in the Appendix, the border would be further pushed down towards
lower-left direction in the figure, enabling us to investigate spectroscopically smaller
region with higher temporal cadence.
In Figure 4.15, two dashed lines from lower-left to upper-right, one roughly
indicating sound transit time scale (with sound velocity set as cs = 100 km/s) and theother Alfven transit time scale, are shown. With the baseline imaging-spectroscopic

4-19

performance of the photon-counting telescope, physical processes with sound time scalein active regions would be observed with angular resolution (size) of 2" and temporal
resolution of ~10 s, while those with Alfven time scale in flares with angular resolution(size) of ~3" and temporal resolution of ~2 s.

4.3

Ultra-High Spatial Resolution Normal Incidence EUV Telescope

As demonstrated by Hinode observations, activities in the lower atmosphere of theSun, such as those of spicules in the chromoshere, and their consequences to the upper
atmosphere, may hold the key to understand energy transfer and subsequent dissipation(heating) into the corona. While dynamics of the transition region and the lower corona
will be extensively studied quantitatively by EUVS, combination with an imaging
telescope with the following features as compared to EUVS should further enhancescience output on the dynamics, energy transfer and dissipation in the transition region
and the low corona.
(1) Temperature overlap(2) Higher spatial resolution(3) Higher exposure cadence(4) Wider field of view
Data from Hinode EIS has clearly indicated that there is structure and velocity
component present whose spatial scale below the instrument resolution. A similar casecould happen with the EUVS instrument under which full interpretation of observed
data, such as line profiles, would confront difficulty. An imaging telescope with theabove-listed features should provide context information for EUVS and would greatly
help to understand what is going on in their data.
As the spatial resolution of the EUVS instrument will be ~0.3" (~0.16" pixel size),that for the imaging telescope should be equal to, or better than, 0.2"-0.3" (~0.1" pixelsize). The need of ultra-high spatial resolution together with temperature overlap with
EUVS leads to the choice of EUV normal incidence optics.
While EUV image data should be quite useful to observe structure, motions andmagnetic environment of coronal phenomena, the image data by themselves would be
not suited to derive physical quantities such as density, actual temperature and velocity
of coronal plasmas. Nevertheless, with the spatial resolution of 0.2"-0.3", the EUV
normal incidence telescope, when combined with EUVS, is expected to contribute in
Solar-C/Plan-B science in the following aspects: (1) Resolve fine, hopefully elemental,
structures and study plasma flows, magneto-acoustic waves and small-scale heatings incoronal temperature range, (2) investigate transfer and dissipation mechanism ofmagnetic energy, and (3) understand how corona reacts to activities that occurred in thelower atmosphere. Furthermore, the movement of the localized particle-acceleration site
can be also investigated by observing the EUV flare kernel with high spatial andtemporal resolution (Figure 4.16).

4-20

Figure 4.16. Target structures and phenomena of the ultra-high spatial resolution normal incidence
EUV telescope.

4.3.1

Scientific Requirements for the Instrument

The primary goal of the ultra high-spatial resolution normal incidence EUV telescopefor Solar-C/Plan-B mission is to give support in understanding physics of magnetic
energy storage and release in coronal plasmas by investigating how corona reacts toactivities that occurred in lower atmosphere. Based on the science goal, scientificrequirements on the EUV telescope are summarized below to deliver largest synergy
with EUVS.
1.

The telescope should have higher angular resolution than EUVS, that is, toresolve structure with 0.2"-0.3" resolution with 0.1" detector pixel for resolving
fine coronal structures to give aid in understanding elemental processes in corona.

2.

Exposure cadence should exceed that for EUVS. As exposure time for EUVS is
1-5 s for active regions with 0.33" spatial sampling while that for the EUV
telescope for active regions and flares can be as short as 1 s and 0.1 s, respectively,
the requirement on cadence can be reasonably met.

3.

Field of view should well cover that for EUVS (that is 200" nominal, >300"
extended). This results in the telescope field of view requirement of ~400"×400".
Note this size covers even outer structures of an active region.

4.

Temperature range should well overlap that for EUVS. This could imply multiple
normal-incidence channels for the telescope. The temperature overlap would also
help in co-aligning image data with EUVS. A UV band for the telescope may also

4-21

be useful for co-alignment with SUVIT data.
5.

To perform observations with low-scattering optics. In order to observe faintstructures in the corona, such as long loops in active regions, plume in coronalholes and structures on the limb, scatter light of the telescope has to be low.

4.3.2

Strawman Instrument Concept

The following depicts preliminary, strawman instrument concept of the EUV
telescope. Specifications of the instrument design should be elaborated in coordination
with the EUVS instrument.
4.3.2.1

Optical design

The instrument will consist of two telescopes. The telescopes have the same opticaldesign except the optical bandpass. The optical design of the telescope is a
Ritchey-Chrétien form of Cassegrain. The primary and secondary mirrors are divided tothree sectors and each sector is coated by the different multi-layer coatings for focusing
in 3 different wavelengths by one optical system.
At the front aperture of each telescope, there are entrance filters and an aperture
selector. The aperture selector conceals two third of the aperture and solar EUV lightincidents only on one sector of the primary mirror. The light then travels the telescopetube to the primary mirror, where it is reflected back toward the secondary. The
secondary is mounted on an active structure that is actuated to remove the instrument
pointing errors and focus the system. Light reflects off the secondary, and toward the
focal plane. There is a shutter and filter wheel near the focal plane. The shutter controls
the exposure time. The filter wheel places the desired filter in the light path, thus
controlling the spectral content of the incoming light. Since the mirror coatings act as
narrow bandpass filters, the focal plane filters effectively select which channel of thetelescope is sending light to the detector. The telescope has the guide telescope fortracking the position of the sun and measuring the instrument pointing error that is used
for image stability compensation. The optical design and the guide telescope system are
similar to that of the AIA/SDO and the TRACE.
4.3.2.2

Exposure time and the diameter of the primary mirror

Based on the refractivity of the mirrors, the transmission of the filters and thequantum efficiency of the CCD used in the AIA/SDO, we expected the effective areas
for EUV lines of the instrument, and predicted the exposure times for accumulating 200
photons/pixel. The threshold count rate "200 photons/pixel" is needed to decrease theerror of the count ratio between the X/EUV bands to less than 10%. The count ratio isuseful for the temperature analysis of corona. Fig 4.17 shows an example of the results.
Recently, there has been remarkable development in CMOS active-pixel image
sensors which appear to be best suited as our photon-counting detector because of theirsmall pixel size (can easily be smaller than 10 μm) and their potential of high frame
readout capability.

4-22

Figure 4.17. The Differential Emission Measures and the predicted exposure time for the EUV lines.
The vertical bar indicates the exposure time for each EUV lines if the geometric area of the mirror foreach EUV line is 200 cm2.

4.3.2.3

Selection of wavelength and temperature coverage

The strawman instrument has 5 narrow EUV bands and one broad UV band. In orderto observe the corona with continues temperature coverage, the five EUV bands arechosen for the instrument based on our knowledge of the solar EUV image taken by the
EIT/SOHO, the TRACE and the AIA/SDO.

Fig 4.18. The temperature coverage of the AIA. (Adopted from the “AIA Concept Study Report”
written by AIA Science Team.)

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Selected band (main EUV line)• 171Å
(Fe IX)• 195Å
(Fe XII)• 211Å
(Fe XIV)• 335Å
(Fe XVI)• 94Å
(Fe XVIII): Temperature range (Log10 T): ~5.9: ~6.1: ~6.3: ~6.5: ~6.9 for a flare (blended with 6.0 component.)

The selected bands include the strong EUV lines and they have already been
observed with 1” spatial resolution by AIA/SDO. The UV band of the instrument issimilar to the UV band of SDO and TRACE. The 1600 Å and 1700 Å continuum
images are taken with the band.
Table 4.5.

Scientific summary of the instrument’s characteristics

Spatial Resolution0.2"-0.3" with 0.1"/pixel

Field of View

~400" × 400"

Wavelength Bands

171 Å, 195 Å, 211 Å, 335 Å, 94 Å and broad-UV band

Exposure time

1 s for an active region
0.1 s for a flare

Cadence

4.3.3< 10 sStrawman design details

Details of the strawman design are shown in Table 4.6.
Table 4.6. The strawman design details of the instrument.

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4.3.4 Technical feasibility for realizing the ultra-high spatial resolution
Many elements of the strawman instrument are flight-proven by AIA/SDO andTRACE. However, the instrument will have 5-6 times better spatial resolution than theexisting telescopes. In order to realize such ultra-high spatial resolution of the
instrument, one should further assess the following issues.
・
・
・
・
・
・
4.3.5

Micro-roughness of the mirrors.
Influence of the boundaries between different multi-layer coatings on each mirror.
Stabilization system for the ultra-high spatial resolution.
Vibration levels of the movable elements, especially those of the aperture
selectors.
Stray lights in the telescope tubes.
Influence of high energy particles in the geo-synchronous orbit.
Further optimization for a better instrument

The optimization of the instrument design is not completed by the NGXT sciencegroup. The science goal and the requirements for the instrument should further be
defined in detail in association with EUVS. Trade-off for wavelength selection should
further be investigated. Number of telescopes, number of wavelength sectors on eachmirror, types of multilayer coating and the diameter of the primary mirrors are closely
related to this issue. The optical design may be further optimized for suitable
observations of the corona in the Solar-C mission.

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Appendix 4.A Event Processing
Since it is highly likely that electrons generated by an incident X-ray photon will
spread over the pixel of incidence, it is indispensable to perform post-detection event
processing to reconstruct correctly energy of incident photons by summing up electrons
distributed over several pixels. Furthermore, since there is possibility that two photons
incident to the same pixel, or to two adjacent pixels, in a single frame resulting in pulse
pile-up (which then results in over-estimated X-ray energies), such pile-up should be
avoided as much as possible with the post-detection event processing. As illustrated in
Figure 4.3, this event processing is to be performed on digitized pixel data, namely,
after analog voltage from each pixel signal is digitized with A/D conversion.
We here assume that all X-ray photons detected by the telescope generate electron
clouds whose size (3-σ diameter in Gaussian profile) less than 3 detector pixels. With
this assumption, events whose signals contained within 3×3 pixel area but not extending
beyond that, such as those shown in Figure 4.19(a) are to be identified as "valid events",
and signals from the central 3×3 pixels collected to reconstruct the energy of an incident
X-ray photon. Note in the right-most panel of Figure 4.19(a) represents a situation that
two-photon incidence cannot be excluded in this particular case hence would result ininclusion of a pile-up event. However, as will be demonstrated later, probability ofoccurrence of such case is sufficiently low.

Figure 4.19(a). Example of events (indicating charge spreading across pixels) that are considered tobe "valid" with 5×5 pixel-area event processing. In each panel, red pixel indicates the X-ray-incident
pixel (5×5 area taken to set the red pixel at its center) while orange pixels indicate pixels with signals
due to charge spreading. The right panel denotes the situation that another X-ray photon was incident(into the dark-purple pixel) in a single frame resulting in a mixture of signal charges in the 3×3 area.
The right panel case cannot be excluded by the event processing.

Figure 4.19(b). Example of events that are to be rejected because of possible contamination by thesecond photon in a frame. Pixel of the second-photon incidence is shown in dark purple with light-bluepixels denote pixels with signals from the charge spreading from the second photon. The middle panel
represents a case where the second photon was incident outside the 5×5 pixel area but some of its
spread charge present in the 5×5 area.

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Figure 4.19(b) shows examples of "invalid events" with possibility that the central
3×3 pixel area is contaminated by charge from the second photon, due to signal
distribution extending beyond the 3×3 area. Hence these events are to be rejected by theevent processing.
Such post-detection event processing has been a standard technique for photoncounting CCD detectors in modern X-ray astronomy satellites such as ASCA and Suzaku(e.g., Koyama et al. 2007), and the same, or very similar, scheme can be applied to ourphoton-counting solar X-ray telescope as well. A big difference here is, however,necessary event processing speed (as determined from level of X-ray flux), is usually
1-2 orders of magnitude higher in the case of solar application than X-ray astronomy; inthe latter case, a typical pixel-processing speed is ~1 Mpixels/s (Dotani, T., private
communication) while the pixel-processing speed necessary for the photon-counting
telescope would be 65 Mpixels/s for photon-counting area of 256×256 pixels and 262
Mpixels/s for 512×512 pixels, with 1000 fps frame readout rate. While this difference iscertainly a technical challenge for realizing photon-counting imagery on the X-ray solar
corona, there has been effort in the X-ray astronomy community to achieve
higher-speed pixel processing for the need of observing bright celestial objects. Anabuki
et al. (2002) reported that the pixel-processing speed of 10 Mpixels/s with 5×5
pixel-area method was achieved with an FPGA by introducing a noble pipeline scheme
for processing series of pixel data. Considering technical progress in FPGAs after 2002,
having pixel-processing speed of a few to ~10 times larger than this value would not be
totally unrealistic. Note parallel event processing using multiple FPGA units is a viable
possibility to reduce processing speed for each unit.
Now let us consider what fraction of the incident photons is detected as "valid
events" by the event processing method with 5×5 pixels. Suppose average number ofphoton incident for each frame sample, for each pixel, is p (photons/pixel/sample). As
the photon incidence events follow Poisson statistics, the probability Ptot of having avalid event (which contains cases such as the right panel of Figure 4.19(a)) is given by
Ptot = (1− e− p ) × e−24 p .
This probability includes cases of, not only pure single event, but also multiple events(pile-up events) whose charge distribution happened to be contained in a single 3×3
pixel area. Since the single-event probability P1 for the central pixel is expressed as
P1 = p × e−25 p ,ratio of single events among detected valid events R1 isp
P
R1 = 1 = p.
Ptot e −1
Figure 4.20 presents Ptot, P1, and R1 as a function of p. Ptot takes the maximum valueof 0.015 at p = 0.041 photons/pixel/sample. This means, with N = 1000 fps frame
readout, 15.0 counts/s/pixel is detected as valid events out of photon incidence rate of
41 counts/s/pixel. Single-event fraction R1 exceeds 90% for p ≤ 0.2 photons/pixel/sample. Figure 4.21 shows valid-event detection rate, in terms of counts/s/pixel (i.e., inevery seconds, instead of every sample) as function of average photon incidence rate to
the detector (in terms of photons/s/pixel) and frame readout rate (fps). Obviously,
fraction of valid events against incident photons decreases as the incident photon rateincreases, because of increased probability of multiple photon incidence in a single

4-27

frame. This is an important aspect of photon-counting observations that excessive
incident photon rate simply results in few number of single-events that are useful for
generating energy spectra. Namely, the incident photon count rate has to be adjusted toan appropriate level so as to have good number of single-events detections.

Figure 4.20. Profiles of valid-event detection probability Ptot (solid curve) and single-event fraction
R1 as function of average photon count rate p (photons/pixel/sample). Scales on the left Y axis apply
to Ptot while those on the right Y axis to R1. Probability of being a single event (P1) is also shown in
the figure as a dotted curve.

Figure 4.21. Valid-event detection rate (in terms of counts/s/pixel) with respect to average photon
count (incidence) rate (photons/s/pixel) and frame readout rate (fps). Profile of Ptot in Figure 6.A-2
corresponds to a horizontal cut at the frame rate of 1000 fps in this figure.

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