Interim Report:

SOLAR-C Plan-B
Toward understanding of the elementary structures& fundamental physical principles that govern
manetized plasma

ISAS/JAXA SOLAR-C Working Groupi

ii

Contents
Executive Summary..............................................................................................................1
Chapter 1: Science Goals and Requirements ................................................................... 1-1
1.1.

Understanding Small-scale Elementary Magnetic Structures .......................... 1-1

1.1.1. Elementary Structures Connecting from Photosphere to Corona .................. 1-1
1.1.2. Elementary Structures of Coronal Magnetic Loops ........................................ 1-2
1.1.3. Elementary Structures in Dynamical Chromosphere ..................................... 1-3
1.1.4. Elementary Structures in Photosphere ........................................................... 1-5
1.2.

Understanding Small-scale Elementary Magnetic Structures .......................... 1-7

1.2.1. Connecting the Chromosphere to the Corona ................................................. 1-7
1.2.2. Magnetohydrodynamic (MHD) Waves in the Solar Atmosphere.................... 1-8
1.2.3. Signatures of Energy Dissipations in the Corona ......................................... 1-13
1.2.4. Behaviors of Magnetic Fields in Corona Explored through Seismology....... 1-14
1.3.

Magnetic Reconnection...................................................................................... 1-15

1.3.1. The Solar Atmosphere as a Laboratory for Magnetic Reconnection ............ 1-15
1.3.2. Magnetic reconnection occurring in a wide variety of plasma conditions .... 1-16
1.3.3. Magnetic reconnection in weakly ionized plasma (solar chromospheres) .... 1-17
1.3.4. Magnetic reconnection in collisonless plasma (corona and transition region).................................................................................................................................. 1-18
1.3.5. Particle Acceleration coupled with Reconnection Dynamics ........................ 1-21
1.3.6. What is the role of 3-dimensionality in magnetic reconnection? .................. 1-23
1.4.

Small-scale Physics Initiating Large-Scale Phenomena .................................. 1-24

1.4.1. Flares and CMEs ........................................................................................... 1-24
1.4.2. Solar Winds .................................................................................................... 1-28
Chapter 2: Solar Ultra-violet Visible and IR Telescope ................................................... 2-1
2.1.

Introduction ......................................................................................................... 2-1

2.1.1. Solar UV-Visible-IR telescope in SOLAR-C Plan-B ........................................ 2-1

iii

2.1.2. Science targets with the instrument ............................................................... 2-2
2.2.

Basic requirements to the instrument ................................................................ 2-4

2.3.

Spectral lines compilation ................................................................................... 2-6

2.3.1. Spectral lines for diagnostics of chromospheric magnetic fields ................... 2-9
2.3.2. Spectral lines for diagnostics of chromospheric dynamics .......................... 2-10
2.4.

Configuration of the instrument ....................................................................... 2-10

2.5.

Optical telescope assembly ................................................................................ 2-12

2.5.1. Optical design ................................................................................................ 2-12
2.5.2. Mirror coating and photon budgets ............................................................... 2-14
2.5.3. Structure design............................................................................................. 2-16
2.5.4. Thermal design .............................................................................................. 2-18
2.6.

Focal plane instruments.................................................................................... 2-21

2.6.1. Interface with the telescope ........................................................................... 2-21
2.6.2. Spectro-polarimeter ....................................................................................... 2-22
2.6.3. Broad-band filtergraph .................................................................................. 2-24
2.6.4. Narrow-band filtergraph ............................................................................... 2-25
2.7.

Image stabilization system ............................................................................... 2-27

2.8.

Contamination control ...................................................................................... 2-28

2.9.

Integration and test plan .................................................................................. 2-28

2.9.1. Integration and initial M1-M2 alignment ..................................................... 2-28
2.9.2. Zero-gravity optical performance test ........................................................... 2-29
2.9.3. Opto-thermal test........................................................................................... 2-29
Appendix A. Diagnostics of dynamics and magnetic fields in the chromosphere ..........a1A.1. Limitations in the Zeeman diagnostics................................................................a1A.2. Prospects of the Hanle diagnostics ......................................................................a1
A.3. Inversion of chromospheric lines .........................................................................a2
A.3.1. Obtaining velocities and plasma parameters ...................................................a3
A.3.2. btaining magnetic fields ....................................................................................a3iv

A.4. Numerical modeling of the chromosphere ...........................................................a5
Chapter 3: EUV/FUV High-Throughput Spectroscopic Telescope .................................. 3-1
3.1.

Performance requirements for the spectrometer ................................................ 3-1

3.1.1. Spatial resolution requirement ....................................................................... 3-2
3.1.2. Temporal resolution requirement.................................................................... 3-3
3.1.3. Broad temperature coverage ........................................................................... 3-4
3.1.4. Necessity of low scattering optics .................................................................... 3-5
3.1.5. Substantial jump expected from the NASA IRIS mission .............................. 3-5
3.2.

Technical feasibility of each key element ........................................................... 3-6

3.2.1. Two element optical layout .............................................................................. 3-6
3.2.2. The telescope mirror and mirror coating......................................................... 3-6
3.2.3. Grating and grating coating ............................................................................ 3-7
3.2.4. Solar blind detectors for EUVS ....................................................................... 3-7
3.2.5. Removal of metal filters and thermal control concept .................................... 3-7
3.3.

Strawman Instrument ........................................................................................ 3-8

3.3.1. Optical Design.................................................................................................. 3-8
3.3.2. Effective Area ................................................................................................... 3-8
3.3.3. Selection of wavelength bands......................................................................... 3-9
3.3.4. Expected count rates and Temperature Coverage .......................................... 3-9
3.3.5. Improvements over previous instrumentation.............................................. 3-11
3.4.

Strawman design details ................................................................................... 3-12

3.5.

Further optimization for better instrument ..................................................... 3-13

Appendix 3-A. EUVS count rates for the brightest lines ........................................... 3-14
Chapter 4:

X-ray Imaging (Spectroscopic) Telescope .................................................... 4-1

4.1.

X-ray Telescope for the Overall Solar-C/Plan-B Science .................................... 4-1

4.2.

Photon-Counting Imaging Spectroscopy X-ray Telescope .................................. 4-2

4.2.1. Scientific Objectives ......................................................................................... 4-2
4.2.2. Required Specification for the Instrument...................................................... 4-6v

4.2.3. Baseline Instrumentation ................................................................................ 4-7
4.2.4. Expected Capability ....................................................................................... 4-13
4.3.

Ultra-High Spatial Resolution Normal Incidence EUV Telescope................... 4-20

4.3.1. Scientific Requirements for the Instrument ................................................. 4-21
4.3.2. Strawman Instrument Concept ..................................................................... 4-22
4.3.3. Strawman design details ............................................................................... 4-24
4.3.4. Technical feasibility for realizing the ultra-high spatial resolution ............. 4-25
4.3.5. Further optimization for a better instrument ............................................... 4-25
Appendix 4.A Event Processing .................................................................................. 4-26
Chapter 5: Spacecraft and Mission Design ...................................................................... 5-1
5.1.

Spacecraft Layout ................................................................................................ 5-1

5.2.

Science Telemetry ............................................................................................... 5-5

5.2.1. Data rate estimate ........................................................................................... 5-5
5.2.2. Data handling on board and data recorder ..................................................... 5-5
5.2.3. Science Telemetry System ............................................................................... 5-6
5.3.

Orbit..................................................................................................................... 5-8

5.4.

Pointing Stability and Attitude Control ........................................................... 5-10

5.4.1. Requirements on the pointing stability ......................................................... 5-10
5.4.2. Attitude control system ................................................................................. 5-13
5.4.3. Flexibility on the role attitude....................................................................... 5-14
5.5.

Spacecraft thermal design to dump heat load from 1.5m telescope ................. 5-14

5.6.

Spacecraft power system ................................................................................... 5-15

5.7.

Science operations ............................................................................................. 5-16

Appendix 5.A Data rate estimate................................................................................ 5-17
References ......................................................................................................................... R-1
Acronyms......................................................................................................................... R-10

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Executive Summary

The Solar-C Plan B mission is proposed to study energy transport and dissipation of themagnetic energy governing the dynamic solar atmosphere through elementary magnetic
structures. For this purpose, the plan B mission has three large state-of-art advanced
telescopes that can perform for the first time high spatial resolution, high throughput, highcadence spectroscopic and polarimetric observations seamlessly covering the entire
atmosphere from photosphere and chromosphere to transition region and corona.
Top-Level Science Goals

At the core of all the questions related to how the Sun floods the heliosphere with hot
plasma, radiation, particles, and magnetic fields is the need to understand the elementary
structures and fundamental physical principles that govern a magnetized plasma.
Theory and laboratory experiments provide many useful insights, but the insights have beenmostly triggered by new observations of the Sun’s atmosphere, which is the only
astrophysical plasma accessible with spatially resolving magnetic structures and dynamics.
By performing high spatial resolution imaging observations under extremely stablecondition never realized so far, “Hinode” opened a new frontier in solar physics researches.
For example, series of high resolution Ca II H images allow us for the first time to identify
waves existing along magnetic fields, which can be considered as a key part in heating of thecorona. Magnetic reconnection, which was firmly settled in solar physics with flareobservations by “Yohkoh”, is now used in explaining many aspects of dynamics observed in
the solar atmosphere. However, details of fundamental physical processes governing theheating and dynamics in the solar atmosphere have not yet been described well with physicalterms. This is because of the lack of spectroscopic measurements, which are most useful todiagnose plasma conditions. With Hinode observations, we also recognize that thechromosphere, which is the interface layer between the photosphere and corona, is the most
key atmospheric layer in understanding the heating and dissipation in the magnetic
atmosphere.
The outer atmosphere of the Sun (or stars) is a complicated system consisting of thephotosphere, chromosphere, and the corona. The energy is built up in the magnetic
atmosphere. The energy is transported via various methods, especially waves along magnetic
fields. The dissipation of the magnetic energy is the central engine for causing variability inthe emissions from the Sun. The energy build up, transport, and dissipation are mixed andcomplicated in the magnetic atmosphere, thus making it difficult to understand the governing
physical processes without spectroscopic and polarimetric measurements over the entire
atmosphere. Thus, the Solar-C Plan B mission is designed to make spectroscopic andpolarimetric measurements over the entire atmosphere with necessary spatial and temporal
resolution.
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Scientific objectives of the Solar-C Plan B mission are summarized in Table 1. As a path
to exploring the fundamental physical principles, we seek to understand:How elementary structures of the magnetic atmosphere responsible for heating anddynamics are created and evolve over each temperature domain of the atmosphere;How the energy that sustains the atmospheric structure is transported through small
elementary structures into large-scale corona and energizes the solar wind; and

How the magnetic energy is dissipated in the astrophysical plasma.

The solar atmosphere consists of elementary structures, one of which is magnetic flux tubes.
The elementary structures shall be understood with their relationship with the magnetic fieldat all relevant temperatures. It is essential to quantitatively trace how they are created andevolve in association with magnetic fields and plasma flows over wide temperature domain.
Magnetic field strength and direction in the chromosphere and corona, which is far above thephotosphere and where the plasma is in low plasma (β<1) condition, is the crucial information
to identify the magnetic structures and their roles for causing heating and dynamical nature in
the atmosphere. Coronal loops are well known magnetic structures in the corona, but EUV
density diagnostics suggests that coronal loops contain elemental structures that are about
10% volume of the coronal loops. The unresolved elementary structures would reflect thesource of energy inputs responsible for creating the hot temperature corona.
The energy for solar activity and heating is transported from the photosphere into thecorona associated with magnetic fields. The origin of the energy is the kinetic energy of thegas convection in and below the photosphere. The magnetic fields rooted on the photosphereplay key roles in transporting the energy upward the chrosmosphere and corona, but we havenot identified the mechanisms because of the complexity of fine structures, although somecandidates have been proposed, such as high-speed upward ejections of plasma, propagationof MHD waves, and winding and braiding of magnetic fields for forming field discontinuities
in the corona. It is extremely important to quantitatively evaluate how much energy is carriedby waves and how the magnetic energy is dissipated in the stratified atmosphere both forheating the chromosphere and the corona and for energizing the solar wind.
A complex system of magnetic fields in the solar atmosphere triggers various types ofdynamics and heating at different timescales and at different temperatures. They require theconversion of magnetic energy to heat plasma, drive flows, and accelerate particle to highenergies. This conversion mechanism is “magnetic reconnection,” which may be the central
engine in many solar phenomena, including the explosion of solar flares, the heating of thesolar corona, and the initiation of coronal mass ejections. The solar atmosphere has a wide
variety of plasma conditions; the inner corona consists of fully ionized, collisionless plasma,
whereas the chromosphere is in weakly ionized plasma condition. We seek to address threeimportant questions about magnetic reconnection. Quantitatively evaluation of fastreconnection in the weakly ionized chromosphere would be helpful to understand what
plasma condition controls the energy release rate. What controls the energy distribution
among thermal, kinetic, and non-thermal energies in early stage of flares? This would help tounderstand a wide variety of solar flares. How can reconnections in small scale develop to alarge-scale flare structure?
Exploring energy transfer and dissipation governing the solar phenomena is directly linked
to the space weather researches. Better understanding and predicting the space weather events
are the basis on human related activities in the space. One of the important goals of solar
physics research is to improve ability to predict geo-effective solar phenomena. High spatial
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Table 1. Solar-C scientific objectives and required instruments
Required Instruments
Section

Scientific Objectives

Specific questions to be answered

SUVIT

EUVS

XIT
PC

NI

A. Elementary
structures of themagnetic atmosphereresponsible forheating anddynamics

1.

What generates dynamical
chromosphere, such as jets?

1.1.3

◎

○

2.

What are the tunnel structures
for energy transport fromchromosphere to corona?

1.1.1
1.1.2
1.2.1

◎

◎

3.

How are waves excited insidemagnetic flux tubes?

1.1.4

◎

○

B. Energy transportand dissipation
through small-scale
magnetic structures1.

How do the observed waves
behave in stratified atmosphereand contribute to thechromospheric heating?

1.2.2

◎

◎

2.

How much energy is carried by
the waves that penetrate to thecorona?

1.2.2

◎

◎

3.

What are signatures of energy
dissipation in coronal loops?

1.2.3

◎

◎

1.

What determines energy
distribution in the energy
conversion by reconnection?

1.3.4
1.3.5

◎

◎

2.

How can reconnections in small
scale develop to a large-scale flarestructure?

1.3.6

○

◎

○

3.

Does fast reconnection reallyoccur in the chromosphere?

1.3.3
1.3.4

◎

○

1.

What are the magnetic fieldstructures during filament
eruptions?

1.4.1

◎

○

○

○

2.

How do magnetic flux and massenergy transport through theatmosphere and initiate flares and
CMEs?

1.4.1

◎

◎

○

○

3.

What types of magnetic
structures at the low atmosphere
are origins of solar winds?

1.4.2
1.1.4
1.2.2

○

◎

C. Magnetic energy
dissipation: Unveil
the mask ofreconnection region

D. Small-scale
physical processes
initiating large-scale
phenomena
regulating space
weather◎○

Note: [Origin] [Transport] [Dissipation]: Color indicates main physical process involvedin each specific question.
Note: SUVIT = Solar UV-Visible-IR Telescope (Section 2)
EUVS = EUV/FUV high-throughput spectrometer (Section 3)
XIT = X-ray Imaging Telescope, PC= Photon counting, NI= Normal incident
3and temporal resolution spectroscopic observations will much improve our knowledge about
the initiation of space weather phenomena. Small-scale physical processes in the emergenceof magnetic flux may play key roles in initiating the phenomena. Also, direct measurement ofmagnetic fields in the prominence or filaments tells how the magnetic field structures ofprominences are evolved and developed to eruptions. Thus, as a goal toward the space
weather researches, we seek to understand:How small-scale physical processes initiate large-scale dynamic phenomena regulating
space weather.
The four science goals defined above are discussed in more details in section 1.

Scientific Measurement Requirements and Proposed Model Payloads

The following scientific measurements is required for the set of mission instruments. All ofthem have not yet realized in any currently operating and planned missions, and they are
unique.
Precise polarimetric (S/N~10-4) and spectroscopic measurements for probing
nature of magnetic fields both in the chromosphere and photosphere.



High throughput spectroscopic measurements for probing nature of dynamics
with high time resolution (0.5-1sec).

Seamless and simultaneous coverage of spectroscopic measurements over eachtemperature domain of the solar atmosphere from chromosphere to the coronaand flares for understanding the entire picture of energy transport and dissipation.

High spatial resolution measurements for resolving elementary structures.

These measurement requirements are realized by the set of the following three science
instruments:
1) Solar UV-Visible-IR telescope (SUVIT)
2) EUV/FUV high throughput spectroscopic telescope (EUVS)
3) X-ray imaging (spectroscopic) telescope (XIT).
EUVS will provide the crucial link between the photospheric and chromospheric magnetic
field and plasma characteristics obtained by the SUVIT and the high temporal and spatial
resolution diagnostics of the corona provided by the XIT. Two possibilities are considered for
XIT: photon counting imaging spectroscopy telescope and ultra-high spatial resolution
telescope.

Solar UV-Visible-IR Telescope
The primary purpose of the UV-Visible-IR telescope (SUVIT hereafter, Table 2) is tophysically capture nature of gas dynamics and magnetic fields in the lower atmosphere, i.e.,from the photosphere through the upper chromosphere, and to understand dynamical
evolution of elementary structures of magnetized atmosphere. The telescope is a4diffraction-limited telescope with a 1.5 m aperture in diameter (the size may be optimizedin the future). It has potential to resolve structures with 0.1 arcsec or better for the firsttime from the space observations (Figure 1), which can reveal dynamical behaviors of the low
atmosphere through elementary magnetic structures and key physics responsible for energy
transfer and dissipation in the photosphere and the chromosphere.
Table 2. Key characteristics of UV-Visible-IR telescope
Item
Telescope
Aperture diameter
Focal plane instruments

Wavelength coverage
Spectral lines(spectroscopic observations)Spectral lines(imaging observations)

Spatial resolution
Exposure time
Spectral resolution
Field-of-view

Description
Aplanatic Gregorian telescope equipped with an image
stabilization system
ϕ1.5 mSpecto-polarimeter (SP)Broadband filtergraph (BF)
Narrowband filtergraph (NF)
280 nm to 1.1 μm(polarimetry in the chromosphere)
He I 1083 nm, Ca II 854 nm(polarimetry in the photosphere)
Si I 1082.7 nm, Fe I 853.8 nm(dynamics in the chromosphere)
Mg II h/k 280 nm(chromosphere)
Mg II h/k 280 nm, Ca II 854 nm, H I α 656 nm,
Na I D 589 nm, He I D 587 nm(photosphere)
Fe I TBD nm<0.1” at UV and visible<0.2” at near-IR(intensity observations) 0.1 - 1 sec(polarimetric observations) 1 – 20 sec~120,000 with the spectro-polarimeter~50,000 with the narrowband filtergraph180” x 180”

SUVIT will take the first attempt to measure magnetic fields in the chromosphere fromthe space. The Hinode Solar Optical Telescope measures magnetic field vectors at the thinphotospheric layer with spectro-polarimetric capability with polarimetric sensitivity of ~10-3,
and scientists have recognized that measuring magnetic fields in the chromosphere,
1,000-2,000 km above the photosphere, is essentially important to explore the stratified
magnetic atmosphere. SUVIT can diagnose magnetic fields both at the chromosphere andphotosphere (Figure 1).
Polarimetric sensitivity of ~10-4 is needed for diagnosing the chromospheric magnetic
fields with combination of the Zeeman effect and the Hanle (scattering polarization) effect.
Comparing with the Hinode Solar Optical Telescope (50cm in diameter), SUVIT canaccumulate roughly one order of magnitude larger number of photons in the same spatialsampling and in the same exposure duration. The polarimetric sensitivity of 10-4 is achieved
with 20 s exposure times for 0.2” pixels.
SUVIT covers a wide wavelength region from UV (~2800 Å) to near infrared (1.1 μm),in which there are several spectral lines best suitable for diagnosing dynamics andmagnetic fields in the chromosphere as well as the photosphere. He I 10830 Å and Ca II
8542 Å spectral lines are identified as best lines for diagnosing magnetic fields and Mg II h/klines at 2800 Å are for diagnosing dynamics in the chromosphere.
5

As the focal plane instruments, two instrument capabilities are equipped to the SUVIT:
Precise spectroscopic and polarimetric measurements allow us to determine physicalparameters such as temperatures, velocities, and magnetic fields. Imaging of intensities andmagnetic fields with high temporal cadence is another capability to capture dynamical
behaviors of magnetic fields within an observing field of view. Narrow-band imager is used
for polarimetric and Doppler measurements at limited number of positions in spectral lines.
Broad-band imager is for best possible high spatial monochromatic images of thechromosphere and photosphere.

Figure 1. Spatial resolution and polarimetric sensitivity achieved with Solar-C UV-Visible-IR
telescope, compared to other space-borne and ground-based large telescope.

EUV/FUV High-Throughput Spectroscopic Telescope (EUVS)The EUV/FUV high throughput spectrometer (EUVS) is a telescope for imaging
spectroscopy in the FUV-EUV region with a resolution and effective area an order ofmagnitude higher than currently available instruments for solar studies (Table 3, Figure 2).
The spatial resolution of EUVS is ~0.3 arcsec (plate scale: 0.16 arcsec/pixel), which is muchhigher than currently operating SUMER and EIS (2-3 arcsec, 1 arcsec/pixel). This higher
spatial resolution is required to spatially resolve elementary structures responsible for energy
transport in the solar magnetic atmosphere. The volumetric filling factor of coronal loops
derived from density sensitive line analysis suggests the existence of fine magnetic structuresin order of 0.3 arcsec or ~200km. Magnetic structures in the chromosphere are also of thissize, as seen as spicules. Thus, this performance allows us for the first time to trace flows of
the energy from its origin (the photosphere) through the chromosphere and transition region
and into the corona. Since the structures are small and dynamic, a high time cadence isnecessary and consequently a high-throughput instrument is required. EUVS will have 1-5 sec
exposure for intense lines with 0.3 arcsec spatial sampling, and 0.5-1 sec or shorter exposurefor 1 arcsec sampling.
Simultaneous spectroscopic measurements sampling all temperature ranges of the solar
atmosphere are essential to achieving the science goals of the Solar-C mission. EUVS is6designed to seamlessly acquire spectral lines emerging from a wide temperature in 104-107
K. Since lower temperature lines are mostly observed in the FUV region and higher
temperature lines are in the EUV region, EUVS is designed to cover a wide wavelength range
in FUV-EUV and it is something like the instrument combining EIS and SUMER.
The strawman spectrometer consists of an off-axis primary mirror and grating mirror. This
optical layout minimizes the number of optical elements, resulting in high throughput
performance. The spectrometer measures emission lines and continua in some numbers ofwavelength windows between ~100 Å and ~1200Å. Low-scattering optics should be used in
the spectrometer to allow us to diagnose low emission measure pixels. Many importantscience discoveries may be hidden in low emission measure regions, such as, reconnection
outflow regions that are much fainter than flare loops, and coronal holes where solar wind
acceleration and shock formation may be observed in the high corona.
Table 3. EUVS scientific characteristic
Item
Instrument design

Plate scale
Simultaneous field of view
Maximum raster width
Wavelength bands
Exposure times
Velocity resolution
Temperature coverage

Description- Φ 30-40 cm, off-axis paraboloid mirror- Single grating spectrograph- Two detectors, one for EUV band and the other for longwavelength bands- Image stabilization, slit-jaw camera~0.16"/pixel
Slit: 0.16", 0.32", 0.96"; Slot: 10", 40", 200"
Slit/slot length: 200" nominal, >300" extended+/- 75" nominal, +/- 200" extended
Four channels (Å):130-250; 510-630; 690-810; 940-1060
1-5 s with 0.33" spatial sampling
0.5-1 s in active regions with 1.0" spatial sampling
Doppler (centroid) shift measurement accuracy <3km/s
Turbulent velocity (width) <10km/s
104 – 107 in active regions with Δ(logT) <0.3

Figure 2. EUVS count rates for active region plasma, compared with currently operating
telescopes (SUMER and EIS) and planned IRIS. At least 100-1200 counts are needed forspectral analysis.
7

X-ray Imaging (Spectroscopic) Telescope
The X-ray imaging telescope images, with spectroscopic information, the uppermost layerof the Sun, the corona, in soft X-rays or EUV to achieve comprehensive understanding on
magneto-hydrodynamic processes across the entire layers of the solar atmosphere jointly with
SUVIT and EUVS. The telescope should aim to (a) understand forms and mechanisms ofstorage and dissipation of energy transferred upwards through the photosphere and thechromosphere, (b) quantitatively understand physics of magnetic reconnection with its
possible role for particle acceleration, and possibly, (c) perform diagnostics on wave
phenomena in the corona. In order to pursue such scientific objectives, the telescope isrequired to have 1) wide coverage on coronal temperatures to observe the entire portion of
single coronal loops, particularly, for the core portion of active regions, 2) high spatial
resolution consistent with that for EUVS, and 3) advanced spectral diagnostic capability forquantitatively investigating physical processes in the corona.
Table 4. Outline of the photon-counting X-ray telescope
Item

Description

Overall description

Perform imaging investigation of the X-ray corona
with the following two modes. Two modes selectable(not simultaneously).
(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 witha segment mirror. A CMOS-APS detector to be used as
the focal-plane imaging array.

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

Mirror focal length

4m

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.)

Field of view

Energy range

Remarks

Telescope plus the focal-planecamera. Includes housings.

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

Data rate* Spatial area size in which a
spectrum with good counting
statistics can be synthesizedwith reasonable integrationtime.

* Off-axis image blur not taken
into account.
** Area with which photoncounting analysis can be
performed.
* Energy range above 2 keVavailable during flares, withphoton-counting mode.

max. 110 Mbps for 512×512 photon-counting area
max. 6.9 Mbps for 128×128 photon-counting area

A photon-counting imaging spectroscopy X-ray telescope (Table 4) is proposed, with aultra-high spatial resolution normal incidence EUV telescope (Table 5) as another option. The
photon-counting X-ray telescope will provide first-ever opportunity to perform
imaging-spectroscopic investigation of the corona in soft X-ray wavelength range, by
measuring energy and location of each incident X-ray photon onto the detector. In addition to8

the photon-counting capability, the telescope is designed to have photon-integration capability
for imaging observations (just like past solar soft X-ray telescopes) with one of the two modes
to be selected for a certain period of observation. The main observational targets for thephoton-counting X-ray telescope is to 1) investigate quantitatively, with spectroscopic
information, structure and dynamics of magnetized plasmas around the site of magnetic
reconnection, and 2) investigate temperature structure of active region corona, in particularthat for high temperatures (above ~5MK). It should be noted that the photon-counting X-ray
telescope provides a unique chance of investigating electron temperatures of thermal plasmas
and even detecting non-thermal emission in soft X-ray wavelengths, which is mostly not
visible with emission lines but should manifest itself as continuum component(s) in X-ray
spectra.
The other option, ultra-high spatial resolution normal incidence EUV telescope, will focus
on understanding connection between the transition region and the lower corona incollaboration with EUVS. The normal incidence telescope serves as delivering
complementary observables with EUVS, providing context imagery information for EUVS
with higher spatial resolution (0.2"-0.3") and wider field-of-view, and with overlapping
temperature coverage. Thus, the main target of this telescope is to resolve fine, hopefully
elemental, structures in the corona, study energy flows and magneto-acoustic waves as well as
small-scale heatings in coronal temperature range.
Table 5. Outline of the ultra-high resolution EUV telescope (Preliminary)
Item

Description

Remarks

Overall description

Consists of two telescopes each with Cassegrain optics.
Each telescope equipped with 3-sector mirrors fordifferent wavelengths.
Image stabilization system necessary.

Optical properties

32 cm aperture diameter. TBD focal length.

Angular 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< 10 sData rate

TBD

TBD

Spacecraft and Mission Design
The spacecraft designed to meet required scientific performance is summarized in Table 6,
and the conceptual layout is shown in Figure 3. The spacecraft is designed mainly withconsidering the following items.
High pointing stability is required to realize the high spatial resolution and precise
polarimetric performance. The spatial resolution of the telescopes is twice to five times
higher than that of telescopes onboard Hinode. For the telescopes with angular resolution
better than about 0.4 arcsec, tip-tilt mirror system is equipped inside the telescopes to9stabilize solar images on their focal plane detectors.
High rate science telemetry is required to acquire spectroscopic and polarimetric datawith high cadence and spatial resolution. The minimum requirement is that the threetelescopes can acquire the data continuously with the post-compression data rate of about
8 Mbps in total. We will use an X-band telemetry system (16-QAM, 16Mbps) with
12-hour or longer duration of contacts at ground stations or a Ka-bad telemetry system(QPSK, ~80Mbps) with shorter duration of contacts at a ground station.
Table 6. Solar-C plan B mission overall characteristics
Spacecraft layout
WeightSize
Launch vehicle
Orbit
Power
Command & Telemetry

Attitude control

Science operations

Telescopes mounted on the optical bench, which is on the Bus
module. See Figure 1.x
4930 kg (at liftoff), 2330 kg (dry weight)
3.2 m x 3.2 m x 7.4 m, excluding solar array paddles
JAXA H-IIA (type 202)
An inclined geo-synchronous orbit (Baseline)
A sun-synchronous polar orbit (optional)~3,500 W (maximum)
Uplink and housekeeping downlink: S-band
Downlink: X-band 16-QAM (16Mbps)
Ka-band QPSK (80Mbps, optional)
3-axis body pointing with high accuracy
Image stabilization system in some telescopes
A function to change the pointing around Z-axis for matching
spectrograph slit direction to the observing target.
Hinode operation scheme, with a final target selection around 10
to 15 minutes before the uplink in cases where this is essential forthe science.

Figure 3. Solar-C plan B mission conceptual layout
An inclined geo-synchronous orbit (GSO, altitude 36,000km, inclination < 30 deg,
period 1 day) is considered as the baseline, with a sun-synchronous polar orbit as the
backup. The GSO orbit provides almost uninterrupted observations of the Sun and muchstable thermal environment, compared with Hinode, for minimizing the thermal
deformation of the structure. It also provides long duration viewing from a groundstation, allowing us to use X-band for high data rate, although the spacecraft weight andsize become bigger.
10

Another advantage of having a long-period downlink is that we can respond to
dynamical changes on the Sun much faster, by monitoring solar images in real time. The
optical and EUV spectrometers have relatively small fields of view and it is important tobe able to choose the observing region as close to real-time as possible.



The spacecraft design should utilize technical heritage from Hinode, especially the
structural design and pointing related technologies for the high pointing stability andtelescope alignment.

11

Chapter 1:

1.1.
1.1.1.

Science Goals and Requirements

Understanding Small-scale Elementary Magnetic Structures
Elementary Structures Connecting from Photosphere to Corona

The Hinode mission provides both the first high-resolution diagnostics of the solar magnetic
field and from a space platform, and is now allowing scientists to discover the processes bywhich the magnetic flux interacts with the solar convective flows in the lower solar atmosphere(the “photosphere”). One of major improvements on our knowledge is that horizontal oriented
magnetic fields are ubiquitously distributed over the solar surface (Lites et al. 2007) and thatthey are dynamically evolved on the photosphere (Centeno et al. 2007, Ishikawa et al. 2007).
Furthermore, Hinode has provided strong indications that the convective motions initiate
outward propagating waves and dynamic reconnection events that result in heating of theplasma accompanied by significant outward mass flux (e.g., Isobe et al. 2008). One of newfeatures revealed with Hinode observations is that many of chromospheric “Type-II” spicules
show substantial transverse displacements during their short lifetimes (<100s), which isinterpreted as propagation of Alfven waves along magnetic field in the chromosphere (DePontieu et al. 2007). This observation also suggests that energy transport takes place via
magnetic structures whose width is ~0.3 arcsec at the chromosphere. Moreover, anotherimportant observational feature is that the spectral lines emitted from the corona attemperatures of 2MK show excess non-thermal line broadening at the footpoints of coronalloops, caused by weak component that is highly blue-shifted to in order of 100 km/s (Hara et al.
2008). Existence of highly blue-shifted features may suggest that numerous number of smalltransient energy releases, i.e., nanoflares, take place at the base of coronal loops (More details
will be discussed in section 1.2).
For understanding the heating of corona and acceleration of the solar wind, it is now the mostessential to quantitatively and exactly identify how the energy (such as mass flow,
magnetohydrodynamic waves with fast mode, slow mode, and Alfven mode, and localized
heating by nanoflares, and flux emergence) is transported to the corona through thechromosphere and transition region, which have a complex system of magnetic fields with fullvariety of plasma dynamics. Observations are required to distinguish the two competingtheories for coronal heating, i.e., mechanical heating by upward-propagating waves, and Joule
heating associated with magnetic reconnections. Quantitative identification is needed for
different magnetic structures, such as active region, quiet Sun and coronal holes, because the
manifestation of the energy transport is varied depending on the natures of magnetic fieldstructures.
To distinguish transport mechanisms, it is very important challenge to resolve the dynamical
nature of small-scale elementary magnetic structures in the complicatedphotosphere-chromosphere-corona system. Furthermore, dynamical nature and evolution ofelementary structures formed in upper atmosphere is directly connected to the formation of
global magnetic field structures and the magnetic energy storage and release that causes theviolent dynamical events (flares and CMEs).

1-11.1.2.

Elementary Structures of Coronal Magnetic Loops

Although the EUV images obtained by TRACE and SDO/AIA (spatial resolution about
0.5arcsec) showed thin coronal loops (Figure 1.1), it is not clear that elemental structures incorona are resolved. Aschwanden & Nightingale (2005) investigated the thermal structure ofcoronal loops based on the TRACE multi-filter analysis, and found that the width of theisothermal loops is 1000~2000 km. If we define the width of an elemental loop by thehomogeneous temperature cross-section, we can say that TRACE resolved the elemental
structure in corona. On the other hand, the EUV spectroscopic measurements with EIS on
Hinode show that the volumetric filling factor of coronal loops is approximately 10% (Warren
et al. 2008); With the spatial resolution of 2-3arcsec (pixel size 1arcsec), only 10% of thevolume in coronal loops are filled with hot plasma. It means that a coronal loop is constructedfrom fine threads with the width in order of ~200 km (0.3arcsec) or less. Thus, there is thedisagreement between the filling-factor analysis based on EUV spectra and the temperature
analysis based on the EUV line images. Doschek et al. (2010) showed that small coronal brightpoints, particularly prominent in coronal holes, has internal structure in the corona, although
their morphological structure are not spatially resolved with the current available instrument(Figure 1.2). Hence, to understand the elementary structures in corona, we need to resolve not
only the stable structures but also the dynamic structures with high time resolution.

Figure 1.1. A magnetically active region at the edge of the Sun shows "loops" that connect one
magnetic polarity to the other. A coronal image (Fe IX in 171A) from TRACE. Coronal
structures seen with 0.5 arcsec spatial resolution.
The EUV spectrometer with the spatial resolution of ~0.3 arcsec, which is about one order ofmagnitude higher than EIS resolution, (and the normal incidence X-ray telescope with thespatial resolution of 0.2 arcsec as the second option), will be able to almost resolve coronalloops at fundamental scales. That is, spatial scales at which loops behave coherently and can be
described by simple hydrodynamic models. The combination of the filling-factor analysis with
the EUV spectrometer and coronal images at 0.1 arcsec pixel from the X-ray imaging telescope
makes clear what the elementary structure of coronal loops is. For this purpose, the X-ray
telescope with normal incidence optics may be more suitable than the photon-counting
telescope.
1-2

Figure 1.2. A polar coronal hole bright point in lines of Fe XI (red), Fe XII (green), and Fe XV(blue). The bright point looks like an unresolved flare seen by Yohkoh.

1.1.3.

Elementary Structures in Dynamical Chromosphere

The chromosphere is the region that interfaces the gas-pressure-dominated photosphere andthe magnetic-field-pressure-dominated corona, through which ionization fraction and theplasma beta change drastically. The chrompsphere forms the conduit for the upward flux ofmass and energy, while most of the non-radiative energy is dissipated there to supply the heat.
Hinode observations demonstrate that the chromosphere is extremely dynamic: small scale
structures called as straws, jets, or fibrils cover the entire surface of the Sun, high speed
chromospheric ‘jets’ (type-II spicules, de Pontieu etal 2007) have a width of ~0.3arcsec andchange with a time scale of 10 second or less. It is obvious that the chromosphere cannot be
considered as an atmospheric ‘layer’, but it is a wholly dynamic 3D entity that never settles tohydrostatic equilibrium. Even in the most quiet area of the Sun, i.e., in the very quiet
photospheric “internetwork”, the chromosphere appears to be dominated by largely acoustic
shocks, which are initiated in the photosphere by the solar convection and the global p-mode
oscillations. Strong “network” fields exist in the photosphere at the boundaries of theinternetwork. These fields spread rapidly with height in the highly stratified chromosphere,
and at some height they form a magnetic canopy over the internetwork regions. Recent
observations from Hinode and other instruments have suggested that the network fields,interacting with convection, energize and structure the quiet chromosphere. Dynamic
phenomena in the chromosphere are more dominating in active regions, where apparent
reconnection events occur frequently, producing outward propagating jets that may be seen atcoronal temperatures (e.g., Shibata et al. 2007, Figure 1.3).
High-speed jets in quiet Sun (type-II spicules) are, however, only visible near the solar limb
with Hinode imaging observations, because of their insufficient optical thickness on the disk
under the Hinode’s Ca II H band-pass filter. In addition, Hinode provides only morphologicalinformation of such chromospheric features. With the lack of the Doppler information, thenature of the type-II spicules in quiet regions or those in sunspot penumbrae (Katsukawa et al.
2007) are still under debate. One of important question is whether they are real plasmaejections with the apparent speed along the magnetic fields or not. Without answering thequestion, we cannot know their physical origin. To understand the generation mechanism andconsequence of this fundamental ingredient of the chromospheres, we definitely need velocity
1-3

information with better contrast and higher spatial and temporal resolutions. The proposed
Solar Ultra-Violet Visible and IR telescope (SUVIT) has a large aperture (1.5 m diameter is
now under study), which can achieve a spatial resolution better than 0.06arcsec in Mg II
280nm line. This UV line provides unprecedented images of fine-scale chromosphericstructures with much higher contrast. The supreme spatial resolution allows us to explore
plasma motions and magnetic-field behaviors inside the chromospheric structures with
~0.3arcsec width. Observations of full spectral line profiles allow us to study the real plasma
flow along the magnetic fields with its acceleration and density distribution, and also toclearly identify the connectivity to their photospheric (or possibly chromospheric) root points.
The high sensitivity, high-resolution Doppler images (Dopplergram) of chromospheres will
for the first time detect the elementary plasma motions, which would be associated with wave
shocks and magnetic reconnections, and thus provide a concrete basis of energetic in thechromosphere.

Figure 1.3. Dynamical chromosphere observed with Hinode. (Left) Dynamical chromosphericfeatures around a large sunspot (Right) Chromospheric jets.
Ubiquitous chromospheric jets associated with cusp-shaped structures (Figure 1.3) show ahint of intermittency of plasma ejection, which infers the presence of more elementary
reconnection processes. SUVIT aims to identify any non-stationary motions or intensity
fluctuation in and around the reconnection site to understand the basic physics that controls
the rate of fast reconnection taking place in the chromospheres (See also Section 1.3). One
particular interest regarding the magnetic reconnection is the injection of magnetic helicity
into the magnetic fields during the reconnection of obliquely directed two flux tubes.
Spectroscopic information in a higher spatial resolution will detect the twist or twisting
motions in a flux tube, which may be crucial importance for dynamics of the ejected plasma
or even for global evolution of the corona. Moreover, SUVIT provides the direct
measurements of the magnetic fields in particular chromospheric features like thereconnection cusp. Even with less spatial and temporal resolution in polarimetric observations,such direct information of magnetic fields at the reconnection site is of a fundamental
importance to establish the reconnection models of the chromospheric dynamics. Enigmatic
phenomena, such as the Ellerman bombs around sunspots (Matsumoto et al. 2008) andintermittent and recurrent occurrence of chromospheric plasma ejections from sunspot lightbridges (Shimizu et al. 2009), are also important targets with a great interest relating to themagnetic field reconnection.
1-4

1.1.4.

Elementary Structures in Photosphere

Photosphere is the region where the gas pressure dominates the magnetic force and, as aresult, the kinetic energy of plasma is transferred into the energy of magnetic fields. The
magnetic fields are discretely distributed and highly localized into tiny magnetic flux tubes.
Hinode SOT aimed to resolve the magnetic flux tubes that have a size of 0.2 - 0.3 arcsec(Figure 1.4, e.g., Nagata et al. 2008), and revealed that they are continuously shaken bysurrounding convective eddies, interacting with each other (merge and split), and occasionally
newly formed and canceled out. These motions of magnetic elements in the photosphere are
the plausible source of the generation of waves propagating upward and magnetic activities
such as ubiquitous reconnections taking place in the upper atmosphere. Individual magnetic
elements are detected, however, with only one or a few pixels in SOT magnetograph (Figure1.4), and it is impossible to infer their internal structures. Thus we are still missing theinformation of electric currents and shape of the cross section of the flux tubes, which are of a
crucial importance for understanding the real energetics of fine scale magnetic fields of theSun. Furthermore SOT frequently observes anomalous Stokes profiles that infer the presenceof unresolved magnetic structures and supersonic flows in photosphere (e.g., Shimizu et al.
2008).

Figure 1.4. A magnetic flux tube captured with Hinode SOT.
The proposed 1.5 m telescope (SUVIT) will observe the interaction between the elementarymagnetic structures and the convection in the photosphere to scrutinize the process of energy
injection into the magnetic fields in much grater detail; with a 0.1arcsec spatial resolution (pixel
scale dx < 0.05”) and with the same photometric accuracy with SOT/Hinode (Figure 1.5),
SUVIT achieves a sensitivity to the magnetic field (dBlong x (dx)2) and the associated energy
elements (j2 = (dBtang x dx)2) one order of magnitude higher than SOT (Figure 1.6). SUVIT
will detect the twist, i.e., the electric current, of individual flux tubes and their evolutions.
Temporal variation of electric current is an important signal of a small-scale energy dissipation
in upper atmosphere. Interaction of flux tubes with vortices in granular convection is ofparticular interest for understanding the generation of electric currents and torsional waves influx tubes. The distribution of magnetic field and velocity across the flux tube and its time
variation will provide a clue to identify the mode of waves propagating along the flux tubes.
Resolving possible ‘hidden’ magnetic polarities in a ‘uni-polar’ region with a higher spatial
resolution is very important to understand the mechanism of ubiquitous chromospheric jets
taking place over there. The most notable advantage of SUVIT is that it can continuously
observe a system of magnetic flux elements in a wide field of view with a highest possible
resolution, accuracy and large time coverage under a seeing-free condition. Such observation
1-5

enable us to track the evolution of foot points of the magnetic field system and, for the first time,
provides a chance to capture a small scale manifestation that triggers the eruption of a magnetic
system.

Figure 1.5. Solar-C, Hinode, and ground-based observations on the plane of the spatial
resolution vs. photometric sensitivity. High S/N observations for chromospheric magnetic field
are shown by open symbols. Detection limits for magnetic flux and energy element are shownby oblique dashed lines.

Figure 1.6. The sensitivity to the magnetic flux and energy elements (εdx2) and the field of view(FOV) of SUVIT and other observations. The SUVIT sensitivity is compared with Hinode SOT
and ground-based observations. The unit in the vertical axis is arbitrary and the horizontal axis
in arcsec. The figure also shows the sensitivity range and FOV of the observations required for
4 specific science topics.
A similar argument can be applied to the Doppler measurements too; i.e., SUVIT has a 10
times higher sensitivity than SOT to the kinetic energy of gas motions. Tiny flows and shocks
possibly associated with photospheric magnetic reconnections, but not resolved so far, are one
1-6

of the important targets of SUVIT observations. Observing the dynamic interactions of themagnetic fields with ubiquitous shocks generated by granular convection in the photospheremay provide a direct clue on the origin of spicules or transient phenomena in chromosphere.
These fundamental plasma processes postulated to be taking place in the high β photosphericplasma are only accessible with high accuracy spectro-polarimetric measurement at 0.1"
resolution under a very stable imaging condition, which can be only realized from the space.

1.2.

Understanding Energy Transport through Small-scale Structures

As briefly described in section 1.1, two non-thermal mechanisms have been debated to
explain the heating of corona and acceleration of the solar wind: upward-propagating waves
along the magnetic structures, and Joule heating associated with magnetic reconnections.
Further observations are required to distinguish how the two mechanisms are excited in the
dynamical atmosphere at the elementary structure scale, how the non-thermal energy can betransported with the mechanisms toward the chromosphere, transition region and corona, andhow the transported energy can be used to heat the atmosphere and accelerating the plasma.
Particularly, the energy transport through small-scale magnetic structures should beobservationally understood to distinguish the roles of the two mechanisms in the heating.
1.2.1.

Connecting the Chromosphere to the Corona

Hinode imaging observations revealed that substantial transverse displacements exist in
many of chromospheric “Type-II” spicules (De Pontieu et al. 2007). Moreover, withcoordination with measurements of EUV spectral lines oriented from the transition region (TR),a blue-red asymmetry was derived in TR spectral lines, suggesting that “Type-II” spicules
transport mass towards the corona (De Pontieu et al. 2009, McIntosh & De Pontieu 2009).
Furthermore, the spectral lines oriented from the corona at temperatures of 2MK show excess
non-thermal line broadening at the footpoints of coronal loops, which can be seen only when
the coronal loops are observed from the top (i.e., disk center view) (Hara et al. 2008). Based onthese Hinode observations (Figure 1.7), quantitatively evaluating of mass and energy transport
between the chromosphere, TR, and the corona at the base of coronal loops is essential to reachthe conclusion for the coronal heating puzzle.
TR (transition region) is the interface connecting between 20,000 K (upper chromosphere)
and 1 MK (corona). SOHO, Hinode and previous space missions have shown the TR to be ahighly dynamic and structured region. SUMER high-resolution spectro-heliograms reveal that,on the quiet Sun, the lower-mid TR (T < 700,000 K) is formed by many loop-like structures 10″
to 20″ long located along and across network boundaries. Similar but fainter structures are also
seen in cell centers. From these data it seems that the majority of the TR emission, particularlyon the quiet Sun, may not represent a continuous transition between the chromosphere and thecorona. EIS raster scans also show the transition from the small-scale TR structures tolarge-scale coronal loops and funnels to occur around 0.7 to 0.8 MK (Matsuzaki et al. 2007).
Of course, transition region “interface” regions must exist at the footpoints of loops or
funnels reaching coronal temperatures but they appear to be responsible for only a small
fraction of the TR emission. This would naturally explain the fact that the TR emits much morethan it would do if it were a very thin (few tens of km) thermal layer, heated by thermal
conduction from the hot corona (as in classical, static, 1-D models). On the other side, SOHO
1-7

observations have also shown the TR to be very dynamic on time scales of the order of 10 s to
100 s.
Quiet Sun TR loops also do not seem to have obvious magnetic counterparts at photosphericlevels, indicating that the loop footpoints are connected to multiple magnetic elements in thephotosphere via a complex system of field lines (Sanchez-Almeida et al. 2007). A similar
phenomenon is also observed in moss areas, regions of enhanced EUV emission in upper TR
and low coronal lines at the footpoints of 3 MK to 5 MK coronal loops in active regions. In thiscase their emission is mainly due to thermal conduction from the hot loops. Each footpoint ofcoronal loops is rooted into a uni-polar magnetic plage region, where no mixed polarity
magnetic flux elements are observed at the photospheric level, suggesting a complex magnetic
coupling in the upper atmosphere (Brooks et al. 2009).
Therefore, measuring magnetic fields in the higher atmosphere, i.e., chromosphere, would be
crucially important for unveiling the nature of this complex coupling between loop-likestructures in TR and the photosphere. At the same time, the EUV spectrometer probes all
temperature ranges from the upper chromosphere up to the corona with high spatial, spectral,
and temporal resolution. The TR strongly emits in the UV/EUV (160 to 2,000 Å) and lines fromincreasingly higher ionization stages generally fall at shorter wavelengths. The EUV
spectrometer should cover a very wide spectral range at unprecedented spatial and temporal
resolution, allowing, for the first time, the full characterization of this region of the solar
atmosphere: from the topology and morphology of the magnetic structures to thethermodynamic state of the plasma (its motions, density, temperature and abundances) that fills
these structures.

Figure 1.7. Schematic illustration of mass and energy transport between the chromosphere, TR,
and corona for the closed loop systems (from De Pontieu et al. 2009)

1.2.2.

Magnetohydrodynamic (MHD) Waves in the Solar Atmosphere

Recent observations, especially Hinode’s imaging observations, have clearly revealed thepresence of apparent transverse oscillations, which are affirmative evidence of the Alfven
waves, in the chromosphere and corona. Wave-like oscillations are observed in a wide variety
1-8

of solar structures, such as prominence in the corona (Okamoto et al. 2007, Figure 1.8), spicules(de Pontieu et al. 2007, Figure 1.9), and Ca II H dynamical jet (Nishizuka et al. 2008) with
Hinode SOT.

Figure 1.8. Example of a prominence thread undergoing synchronous oscillation along its
entire length (from Okamoto et al. 2007)

Figure 1.9. Example of the transverse displacement of a spicule (from de Pontieu et al. 2007)
However, we did not identify whether the observed waves are really propagating Alfven
waves, or just standing waves in most of observations, except for ubiquitous upward waves in
the corona detected with a coronagraph measurement of line-of-sight velocity fluctuations(Tomczyk et al. 2007). Thus, in order to evaluate the propagating MHD waves and their modes,
observations require not only velocity and photometric measurements (Kitagawa et al 2010)
1-9

but also the measurements of the magnetic field fluctuation, as examined with photosphericmagnetic-field data by Fujimura and Tsuneta (2009).
1.2.2.1. Relevance of Observed Waves for Chromospheric Energy Budget
The chromosphere’s temperature is approximately 10,000 K with high density. To maintainthe temperature structure of the chromosphere, it is required to input heat in the chromosphereabout one order of magnitude higher than that required for heating the corona.
Recent studies show that there is insufficient acoustic wave power propagating at
higher-frequencies (>5mHz) from the solar interior to heat the chromosphere (Fossum &Carlsson 2005; Carlsson et al. 2007). However, it is shown that inclined magnetic field lines atthe boundaries of large-scale convective cells provide portals through which low-frequency(<5mHz) magnetoacoustic waves can propagate in the chromosphere (Suematsu 1990), andthese waves provide a significant source of the energy necessary for balancing the radiative
losses of the ambient chromosphere (Jefferies et al. 2006). When magnetoacoustic waves
propagate upward in strongly stratified atmosphere such as the chromosphere, the waves will
steep into shocks to heat the chromosphere (Carlsson & Stein 1997). Carlsson, Hansteen &Gudiksen (2010) indicates that as low-frequency magneto-acoustic shock waves interact withsmall magnetic elements, joule heating may be effective at these elements’ edges.
It is important to observationally determine how much energy is sustained in observed waves,with identifying the mode of the waves. The energy should be measured as a function ofmagnetic field inclination and atmospheric height. To determine the power of magnetoacoustic
waves at the lower atmosphere, we need the measurements in cadences of better than 10-20
seconds with line-of-sight velocity amplitudes of ~0.1 km/s for the photospheric lines and with
~1 km/s for chromospheric lines (e.g., Vecchio et al. 2007). High spatial resolution, 0.1-0.2
arcsec, is needed to resolve individual channels for wave propagation (Bogdan et al. 2003).
Moreover, determining the importance of shocks is vital for balancing the radiative losses in the
more magnetized parts of the upper chromosphere. To do this, velocity measurements with UVlines in formation temperatures between 104 and 105 K need to be coordinated with thephotospheric and chromospheric measurements with visible lines. For measurements with UVlines, we need cadences of better than 20-30 seconds, which provide 6-15 measurements in one
period of 3-5 minutes oscillations, and Doppler velocity measurements of better than 10 km/s
(equivalent to sound speed in the upper chromosphere) with line width measurements of better
than 5 km/s. It is also important to have the capability to resolve the elemental structures where
the wave power is channeled. These are typically on the order of 0.3-0.4” in size, inferred fromthe observed width of spicules. The wide temperature coverage will allow us to track theevolution of the wave power as it propagates through the different layers of the solar
atmosphere. We should note that the NASA IRIS mission, which will be launched in late 2012,
will provide new information on the chromosphere, especially on the line-of-sight velocity
amplitude measured with chromospheric lines, but it does not have adequate transition region
and coronal coverage, nor will it measure chromospheric magnetic fields, both of them areessential to evaluate the natures of magnetoacoustic waves in the solar atmosphere.
1.2.2.2. Time Series Spectra at Different Atmospheric Layers
The problems encountered in understanding the coupling between the chromosphere and thecorona are illustrated in Figure 2.2-4. The chromosphere has relatively large patches (10′′ to20′′) of many thin, bright spicules (width~0.3”), and the coronal structures are narrow, isolated
loops that fade at their footpoints. There is little correspondence between the two. A
1-10

comparable transition region image that can only be obtained by a rastering spectrometer would
take almost two hours to complete with the current available spectrometer such as SUMER.
Although the large-scale structures do not normally change on this timescale, the small-scale
dynamics are much faster, and thus high-cadence spectroscopic measurements in transition
region lines will make a giant step forward. Moreover, matching spatial resolutions both below(chromosphere and photosphere) and above (transition region and corona) is crucially
important to trace the energy flow over the entire atmosphere.
As illustrated in Figure 1.10, spectral lines show different characteristics among
chromosphere (Cauzzi et al. 2007), transition region, and corona. The chromospheric spectra
have characteristic 3 to 5 min periods, the transition region brightenings occur in bursts of up to
30 min (Innes et al. 1997), while coronal variations are more diffuse. Such different
characteristics would reflect the energy transport and release at each layer, but no one knows
details, because this kind of time series spectra has never been recorded simultaneously withdifferent spectral lines covering the solar atmosphere with 0.3” spatial resolution.

Figure 1.10. The chromosphere-corona connection. The top panels show cut-outs taken fromthe chromosphere (Ca II, with spatial resolution ~ 0.3”) and the lower corona (171A, withspatial resolution ~1′′) images. The regions are outlined in yellow. The bottom right panel
shows representative time series of chromospheric (Ca II), transition region (Si IV), and lower
coronal (Ne VIII) line profiles (recorded at different regions at different times).
The diffuse spectra in coronal lines is much wider than thermal line broadening due to hightemperature plasma, and the extra broadening has been considered as non-thermal line
broadening. The origin of the non-thermal broadening is uncertain, but it may reflect turbulent
nature in the coronal loops. Hinode observations have revealed that the non-thermal line
broadening is well observed at the footpoints of coronal loops and that the non-thermal
broadening is more enhanced when the coronal loops are viewed from the top, suggesting that
there are high-speed upflows which are not resolved spatially (Hara et al. 2008). With having
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shorter exposure with ~0.3” spatial resolution, the non-thermal broadening may be partiallyresolved as Doppler shift signatures of jets in each elementary magnetic structure.
1.2.2.3. Alfven Waves’ Contribution to Heating Corona and Accelerating Solar Winds
Alfven waves, transverse magnetohydrodynamic waves that can propagate along themagnetic field lines over large distances and transport magneto-convective energy from thephotosphere into the corona, are one promising candidate for the energy source both for theacceleration of the solar wind and for the heating of the quiet corona.
Numerical simulations have been made these days to investigate the nonlinear propagation of
Alfven waves in the highly stratified atmosphere. Alfven waves are significantly attenuated in
the stratified atmosphere between the photosphere and transition region by various dissipation
processes such as nonlinear mode conversion, Alfven wave resonance and phase mixing.
Suzuki and Inutsuka (2005) showed that the limited portion of the waves is able to propagate to
the corona and nonlinear low-frequency Alfven waves can simultaneously heat the corona anddrive the fast solar wind. By using the temporal spectrum of the convection motions at thephotosphere from Hinode observations, Matsumoto and Shibata (2010) showed that theobserved convection motions generate Alfven waves and that the region between thephotosphere and the transition region becomes an Alfven resonant cavity.
The currently available observations of oscillation events do not allow us to determine
whether sufficient energy would remain in the waves that penetrate all the way to the corona tomake a substantial contribution to the heating budget of the solar corona and/or the accelerationof the solar wind. For allowing this evaluation, observations are required to clearly identify
propagating Alfven waves and their modes in the region between the photosphere and thetransition region. Such observations are the measurements of the magnetic field fluctuation at
the chromosphere and the simultaneous measurements of velocity and photometric informationfrom the chromosphere to the transition region.
1.2.2.4. Specific Topics to be Addressed by New Observations
In addition to evaluating how much energy is sustained in observed waves over theatmosphere, the following questions can be newly addressed with these unique observations.
All of them cannot be studied at all with currently available instruments. First, mode conversionof waves may occur around the layer where the plasma beta becomes unity (e.g., Bogdan et al.
2003). Since the beta~1 layer is around or below the middle chromosphere, we may detect thesignature of mode conversion in waves. The phase relation between velocity and magnetic field
fluctuations will tell how observed waves behave in stratified atmosphere, depending on
magnetic structures. Secondly, high throughput performance of the proposed EUV/FUV
spectroscopic telescope can allow a search for higher frequency waves and a search for wave
signatures in line widths and velocity in the transition region and corona. For active region
observations, the telescope will have 1-5 sec exposure time for intense spectral lines with 0.3
arcsec spatial sampling, and 0.5-1 sec exposure time for about 1 arcsec spatial sampling.
Slit-and-stare measurements will provide the temporal series of spectral lines with higher
cadence than 1 sec, allowing a search for higher frequency waves and thus obtaining new
insights in heating of the active region corona by high frequency waves. Thirdly, we may obtainfine-scale structure of magnetic fields by measurements of vector magnetic fields at thechromosphere. We can directly study the properties of turbulent cascade or phase mixing of
Alfvenic waves by these small scale magnetic structure. For instance small-scale current sheets
1-12

in the chromosphere would be an evidence for phase mixing. Anisotropic elongated structure
along magnetic fields would be an evidence of cascade of Alfvenic turbulence.

1.2.3. Signatures of Energy Dissipations in the Corona
An outstanding issue that remains elusive in the formation and evolution of coronal loops isthe localization of energy conversion. As an alternative to the classic nanoflare model withmagnetic reconnection at coronal heights (Parker et al. 1989), it has been recently suggested
that reconnection takes place in the chromosphere and is responsible for the coronal heating
through the formation of jets and spicules (e.g., Aschwanden, Nightingale & Alexander 2000).
As mentioned previously, recent Hinode observations indicate high-speed upward mass flows
at small scale and excess non-thermal line broadening at the base of coronal loops, suggestingthat the origin of coronal heating exists in the base of the coronal structures (section 1.2.1).

Figure 1.11. Left: Differential emission measure (DEM) distribution for an active
region. Note no emission line is available for temperature diagnostic beyond logT = 6.7 (T= 5 MK). Right: DEM distribution for quiet Sun. After Brosius et al. 1996.

High temporal and spatial resolutions are needed to resolve the evolution of numerous,
small-scale explosive events like a nanoflare “storm” (Klimchuk, 2006). This tells us whether
continuous heating changing slowly compared to a cooling time or impulsive heating repeating
rapidly compared to a cooling time is more significant for the coronal heating process. In
addition, plasma diagnostics capabilities in the broad temperature coverage, such asmeasurements of the differential emission measure (DEM, Figure 1.11), electron density, or
turbulent velocities, are crucial to advancing the understanding of coronal heating. Also, thepair of EUV emission lines from different atomic species allows us to estimate the thermal
temperature of ions because of the difference in their thermal velocities (Imada et al. 2009). The
EUV spectrometer proposed for Solar-C is the ideal instrument for the plasma diagnostic
because of its sub-arcsec spatial resolution, EIS-like or better spectral resolution, wide
temperature coverage from the photosphere into the corona, and high throughput. The EUV
spectrometer has the diagnostic capabilities for the plasma only up to ~5 MK, because
high-temperature plasma in 3-6MK is energetically dominant in active-region corona (Figure1.11). There is also an interesting indication from Hinode XRT that there possibly exist
high-temperature components whose temperature at least 5 MK (could be as high as 10 MK)
diffusely distributing across the quiet Sun, even without any active regions (Ishibashi et al.
2011). High-temperature components can be created by energy dissipations in the corona.
1-13

Identification of temperatures of such components in the corona, coupled with their spatial
distribution, may provide important clues on how the corona is heated, and maintains its
temperature. If the proposed photon-counting X-ray telescope is selected as the X-ray imaging
telescope, it will reveal high-temperature temperature components and their temporal evolution
by obtaining spatially-resolved X-ray count spectra.

1.2.4. Behaviors of Magnetic Fields in Corona Explored through Seismology
It has been difficult to measure magnetic field fluctuation in the corona, which is one of
key information to investigate propagating waves. However, observations of oscillations in
the corona have the potential to reveal different properties of the local plasma environment. For
example, kink-mode transverse oscillations of the coronal loops impulsively excited by theblast wave in a large flare (Nakariakov et al. 1999, Figure 1.12) can be recognized in the timeseries of coronal images and allow us to estimate the value of the local magnetic field strength.
In some cases, both fundamental mode and higher harmonics are obtained, which enables us toestimate the density stratification along coronal loops.

Figure 1.12. Temporal evolution of the loop displacement through the loop apex whenoscillation is exited in the loop (Nakariakov et al. 1999)
The Solar-C observations will allow us to take Doppler velocity measurements of such
oscillating loops. Oscillations of coronal loops during Alfven waves propagate along theloops may be measured with the proposed EUV spectrometer, which may derive local
magnetic field strength as a function of the distance from the base of coronal loops. The
techniques of coronal seismology require cadences of the order of 2 s (Roberts 2000) to
determine radially dependent parameters. By having higher spatial resolution (~0.3”), thedensity and temperature can be measured in the configuration that different coronal loops areless overlapped. This allows to have better estimate on the density and the local field strength.
Direct measurement of velocity is a key parameter, because linear Alfvén waves areincompressible and so do not show up as intensity perturbations. The velocity resolution
needed to observe the perturbations caused by these waves is of order 10 km s-1 to see spatiallyresolved waves and 5 km s-1 should be these waves on smaller scales than we can resolve.

1-14

1.3.
1.3.1.

Magnetic Reconnection
The Solar Atmosphere as a Laboratory for Magnetic Reconnection

Magnetic reconnection is one of the most important mechanisms for the plasma heating andparticle acceleration in magnetized space and astrophysical plasmas. The importance ofmagnetic reconnection is well discussed in many plasma environment, such as the Earth’s
magnetosphere (e.g., Hones, 1979; Nagai et al., 1998, 2001; Baumjohann et al., 1999; Øieroset
et al., 2002; Imada et al., 2005, 2007, 2008), laboratory (e.g., Baum & Bratenahl, 1974; Ono etal., 1997; Ji et al., 1998; Yamada et al., 1997), and also the solar atmosphere. Several
astronomical phenomena, such as accretion disks, proto-type stars, AGN jets, cluster ofgalaxies, and pulser wind, are also recently simulated and discussed with magnetic
reconnection to explain dynamical natures revealed with X-ray and radio observations.

Figure 1.13. Global magnetic configuration and associated reconnection signatures in a type ofsolar flares (Tsuneta et al. 1997).
Magnetic reconnection is believed to be central to the energy release in solar flares, theheating of the solar corona, and the initiation of coronal mass ejections (CMEs) (Figure 1.13). It
is a rapid mechanism for converting the stored magnetic energy into kinetic, thermal,
non-thermal and wave/turbulence energy via magnetic dissipation process. Yohkoh X-ray
observations provided some pieces of evidence for magnetic reconnection in solar flares andcoronal jets, and then Hinode has revealed that magnetic reconnection occurs with various
configurations of magnetic fields not only in the corona but also in the chromosphere. Theenergy conversion by magnetic reconnection is fundamental and essential to explain dynamical
behaviors in the wide variety of plasmas in the solar atmosphere. While considerable progresshas been made over the past decade, three key questions listed below have not yet beenanswered with observations:

1-151) What plasma conditions and parameters, such as collisionality or temperature, control theenergy release rate of magnetic reconnection (reconnection rate)?
→ Section 1.3.2 and 1.3.3
2) What determines the energy distribution among kinetic, thermal, wave/turbulence, andespecially non-thermal energy at the rapid conversion of stored magnetic energy?
→
Section 1.3.4 and 1.3.5
3) What is the role of 3-dimensionality in magnetic reconnection?
→ Section 1.3.6
Observations carried out by the Solar-C Plan B mission should be designed to answer thesequestions.

1.3.2.

Magnetic reconnection occurring in a wide variety of plasma conditions

One of the fundamental questions in magnetic reconnection is what determines the energy
release rate. It is also what fraction of released energy is converted into kinetic, thermal,
non-thermal, and wave/turbulence energies. Previous studies of solar, space and laboratory
plasmas challenged these questions, though they have not yet answered these questions. The
solar atmosphere shows us a wide variety of plasma conditions, i.e. weakly ionized to fully
ionized plasmas, and collisional to collisionless plasmas (Table 1.1), which are quite different
from the plasma conditions in magnetosphere or laboratory experiments. Observational
investigation of magnetic reconnection under various plasma conditions is very important forspace, solar, astrophysical and laboratory plasma physics. The challenge of the Solar-C mission
is to diagnose magnetic reconnection in such various plasma conditions with more direct
observational information available from spectroscopic measurements in three different
wavelengths (IR-visible-UV, FUV-EUV, and soft X-ray). Especially, the proposed highthroughput UV/EUV spectroscopic telescope can reveal plasma dynamics in faint reconnection
region with high cadence measurements, and the proposed photon-counting X-ray telescope
will enable us to diagnose particle acceleration mechanism with imaging spectroscopy in
X-rays. The spectro-polarimetric instrument of the proposed IR-visible-UV telescope measuresthe magnetic fields in the chromosphere, allowing us to understand chromosphericreconnection more detail.
Table 1.1. Typical properties of plasma in solar atmosphere
Photosphere
Dominant pressure
Ionization
Collisionality

Gas
NeutralCollisional

Chromosphere

Inner Corona

Magnetic field
Weakly ionized
Collisional

Magnetic field
Fully ionizedSemi-collisionless

1-16

Outer Corona(~2Rs)
Gas
Fully ionized
Collisionless1.3.3.

Magnetic reconnection in weakly ionized plasma (solar chromosphere)

1.3.3.1. Does fast reconnection really occurs in the chromospheres?
One of the most important Hinode findings is that magnetic reconnection events frequently
occur in the solar chromosphere. Shibata et al. (2007) found the ubiquitous presence ofchromospheric anemone jets outside sunspots in active regions. The anemone jets have an
inverted Y-shape, similar to the shape of X-ray anemone jets in the corona (Figure 1.3). Theirtypical velocity is comparable to the local Alfven speed in the lower chromospheres, and theirfootpoints are located among the mixed polarity region. So they concluded that chromosphericanemone jets imply magnetic reconnection in the chromospheres.
Hinode has also discovered new types of magnetic reconnection events in the chromosphereswith various magnetic field configurations (e.g. Katsukawa et al. 2007, Nishizuka et al. 2008,
Shimizu et al. 2009). Shimizu et al. (2009) revealed that sunspot light bridges sometimes show
recurrent occurrence of chromospheric jets. Vector magnetic field data in the photospheresuggests that a helical magnetic flux tube lying below the light bridge plays vital roles in
producing jets. Katsukawa et al. (2007) also found microjets even in the sunspot penumbra,
where magnetic field lines are usually not fully anti-parallel.
One of the important questions about magnetic reconnection in the solar chromosphere iswhether the magnetic reconnection in the chromosphere is really fast, independent on
resistivity, or not. Our current guess, which is based on Hinode observations, is ”yes”. Hinode
Ca II H movies, however, do not allow us to evaluate the reconnection rate. The time cadence of
the Ca II H observations is fast enough, but quantitative information, such as plasma flows,density, temperature and magnetic field, are needed. These quantitative parameters areavailable from the proposed IR-visible-UV telescope
1.3.3.2. Observations of Ambipolar effect in weakly ionized plasma
Chromosphere is high density, collisional and weakly ionized plasma. In such a condition, it
is hard to expect fast reconnection. On the basis of the space and laboratory plasma
observations, it is said that, to drive “fast reconnection”, “anomalous resistivity” should work in
the diffusion region via microscopic process. To make it possible, the following two aspects
should be satisfied: ion gyro radius is larger than mean free path (collisionless condition), andthe width of a current sheet becomes thin comparable to ion gyro radius. However, it seems that
Sweet-Parker current sheet is not so thin, and the mean free path in the lower chromosphere issmaller than ion gyro radius. If fast reconnection really occurs in the chromosphere, additional
physics should be considered.
One candidate is “ambipolar effect” in the chromosphere, where the neutral hydrogen to
proton density ratio is ~104 (i.e. Singh and Krishan 2010). Ambipolar diffusivity is effective inthe chromosphere because of the ion-neutral collisions. The collision between ions and neutrals
causes ambipolar diffusion, which plays a role in making a current sheet thinner and thinner inthe weakly ionized plasma (e.g., Zweibel, 1988, 1989, Brandenburg & Zweibel, 1994, Chiueh,1998). Isobe et al. (in preparation) discussed the effect of ambipolar diffusivity and showed thatbursty reconnection can take place in the chromosphere (Fig 1.14). The nature of magnetic
reconnection in the chromosphere is still not clear from both observational and theoretical
viewpoints. The aim of the Solar-C mission with the IR-visible-UV spectro-polarimetric1-17

observations is to measure physical parameters around the reconnection region andobservationally determine the reconnection rate qualitatively and to reveal the dynamics ofmagnetized plasma in various configurations of magnetic field with direct observation ofchromospheric vector magnetic field.

Figure 1.14. Magnetic reconnection in weakly ionized plasma (bottom) and collisionlessplasma (top).
1.3.4.

Magnetic reconnection in collisonless plasma (corona and transition region)

Magnetic reconnection in the solar corona has been discussed as a key mechanism involved
in solar flares. Various features expected from the magnetic-reconnection model have been
confirmed with modern space observations. For example, Yohkoh observed the cusp-likestructure, which is predicted by the standard reconnection flare model (Tsuneta et al., 1997).
Quantitative understanding in the key region of solar flares, however, is not enough so far to
discuss the energy release rate or energy distribution rate of magnetic reconnection.
The primary reason why it is difficult to directly observe the magnetic reconnection region is
because of its faintness. A bright cusp-like structure (flare loop) is formed during a solar flare,
and the reconnection region that is believed to exist above the cusp structure is extremely faint.
A large dynamic range of intensity and low-scattering performance is required to observe inside
the reconnection region. Moreover, since the temperature in the reconnection region may besignificantly higher than that of cusp-like flare loops (e.g., Tsuneta et al., 1997), the line
spectral signature (ionization state) in the reconnection region may be different from what seenin flare loops.
1.3.4.1. What determines the energy distribution
wave/turbulence, and non-thermal energies?

amongkinetic,

thermal,

One of the important parameters for magnetic reconnection in the solar corona is how muchenergy is stored in the coronal magnetic fields. The stored energy can be estimated from themagnetic-field measurements at the photosphere (Top part of Fig 1.15). For example, Kubo etal. (2007) showed magnetic-field evolution of a flare productive active region with continuous
series of photospheric magnetic fields by Hinode. Magara and Tsuneta (2008) calculated themagnetic helicity in the flare-productive active region and found that the helicity increases very
rapidly before the flare. This quantitative analysis is very important for the estimation of stored
energy before magnetic reconnection. Reconnection electric fields (reconnection rate) also can
1-18

be quantitatively estimated with Hinode (e.g. Jing et al., 2008). Therefore, we canobservationally evaluate the stored magnetic field energy and its release rate, which is the top
part of Fig 1.15.
To understand the energy distribution among the energies, we need to evaluate each ofthermal, kinetic, non-thermal, and wave/turbulence energies during a flare. Spectroscopic
observations with EIS on board Hinode enable us to evaluate thermal energy with differential
emission measure analysis (e.g., Watanabe et al. 2007), kinetic energy with the Doppler
velocity analysis (e.g., Asai et al., 2008), wave/turbulence energy with line broadening analysis(Imada et al. 2008), and non-thermal energy with simultaneous RHESSI hard X-ray andmicrowave observations (e.g. Minoshima et al. 2009, Watanabe et al. 2010).
As shown in Figure 1.15, with Hinode observations, we can know the magnetic energy stored
in corona before magnetic reconnection (top part of Figure 1.15). The most part of energies canbe evaluated in the post reconnection stage after the transient energy release is finished (bottom
part of Figure 1.15). On the other hand, there is poor observational knowledge on the physicalparameters in the reconnection region during the period when the reconnection process is on
going (middle part of Figure 1.15). The inflow into the reconnection region, the plasma
temperature in the reconnection region, and the temperatures and densities of the plasma jets
(outflows) have not been measured quantitatively. They are key observational parameters
necessary for probing reconnection process.

Figure 1.15. Energy distributions during magnetic reconnection.
1.3.4.2. Probing inside the reconnection region in non-ionization equilibrium
Theoretical model of steady fast reconnection, called Petchek-type reconnection, expected
the formation of two slow-mode shocks elongated from the reconnection region, which play arole in heating ambient plasma. During rapid heating with slow-mode shocks, heated plasma
would be in non-thermal equilibrium condition. There are typically three kind of thermal
non-equilibrium conditions; 1) Non Maxwellian energy distribution, 2) different temperature indifferent species (e.g., Tp > Te), and 3) transient ionization condition (ionization
1-19

non-equilibrium). All of three conditions are predicted in the reconnection region, and thedetection of them are a challenging target for Solar-C.
Figure 1.16 shows a calculation of non-equilibrium ionization structure in 1200 km/s jets atthe down stream of magnetic reconnection region. We can clearly see that the plasma state in
the downstream of slow shocks cannot reach to ionization equilibrium. The proposed
high-throughput EUV spectrometer (~10 times better than the effective area of the currently
available spectrometer, see Figure 3.3 in Chapter 3) can measure spectral lines with much
shorter exposures, allowing detecting the transient state of non-ionization equilibrium. Suchtransient state tells us a history of ionization process of each ion and thus contains theinformation of the reconnection region. For detecting the transient ionization condition, a lot ofspectral lines covering a broad range of ionization states need to be measured, which can bedone only with the proposed UV/EUV spectrometer (for flare plasma, from Fe XVIII to Fe
XXIV, see Appendix 3-A).

Figure 1.16. A computation of non-equilibrium ionization structure in 1200 km/s jets at the
down steam of magnetic reconnection region. The following parameters are used in thiscomputation: magnetic field 12 gauss, electron density 4.2x108cm-3, temperature 1MK, andplasma β 0.01 at the upstream side of the slow-mode shock and temperature 40MK at thedownstream; the shock angle at the upstream and downstream is 85 deg and 2.3 deg,
respectively. The upper right panel shows the total intensity of the non-equilibrium ionization
structure (black line) as a function of ionization state is compared with the total intensity of
post-flare loop (blue dotted line).
The signal intensity of the non-equilibrium ionization structure can be seen in contrast withthe bright post-flare loop, if we choose suitable spectral lines (Fe XV ~ Fe XVIII). Bright flare
1-20

loops have typical temperature of ~10 MK and density of ~1011cm-3, and therefore the loops are
dominantly observed with Fe XIX ~ Fe XXII. In the reconnection region, it takes ~100s toreach the ion state to higher ionization states after the energy release in a short time scale. Thus,
we expect that signals exist not only in high ionization states (Fe XVIII ~ Fe XXIV) but also in
low ionization states (Fe XII ~ Fe XVII).
When spectral lines with ionization states higher than Fe XIX are used for measurements,
observing the early phase of solar flares is a key to success the diagnostics of the reconnection
region. This is simply because bright flare loops developed during flares has higher density, say,by two orders of magnitude, than the density inside the reconnection region. This means thatthe spectral line intensities from the flare loops are higher by four orders of magnitude than thatof the reconnection region. In the early phase, the total emission measure of flare loops is still
small and thus we have possibility to detect signals from the reconnection region, when we
have the spectrometer with wide dynamic range and low scattered-light optics.
In the down stream of slow-mode shocks, the ion temperature may not be same as the electron
temperature. Such non-equilibrium is frequently observed in collisionless plasma condition(Baumjohann et al., 1999). To understand how much energy is given to ions and electrons
respectively, we should observe both of ion and electron temperatures simultaneously. Linewidth of EUV spectral lines can be used to estimate ion temperature (e.g., Imada et al., 2009),
but the ionization equilibrium is always assumed in the spectroscopic analysis of EUV lines,
resulting in difficulty in deriving electron temperatures for transient non-equilibrium plasma.
Therefore it is essential to have another measurement for deriving the electron temperatureinside the magnetic reconnection region. Accurate measurements of the electron temperature
can be achieved by measuring thermal bremsstrahlung emission at short wavelengths (soft
X-rays) with the proposed photon-counting X-ray telescope. In contrast to the EUV line
spectroscopy, the spectral shape of thermal bremsstrahlung emission directly reflects theelectron temperature without depending on the ionization equilibrium.
Because of its faint structure, non-ionization equilibrium signals inside the magnetic
reconnection region are not easy to detect. But the proposed UV/EUV high-throughput
spectroscopic telescope has potential to detect them, because of its high-throughput
performance and low scattered optics. If we can observe inside the reconnection region, thenwe can determine the reconnection condition by using the knowledge of non-ionization
equilibrium. The ionization process is similar to the nuclear-decay process. Thus we canestimate the time passage from strong heating. The velocity, density, and ion temperature in thereconnection region can be derived from line spectroscopic analysis with the knowledge of
non-ionization equilibrium plasma. Therefore, probing the magnetic reconnection region fromviewpoint of non-ionization equilibrium can provide useful information on the mechanisms of
the energy distribution rate.
1.3.5.

Particle Acceleration coupled with Reconnection Dynamics

Yohkoh observations played an important role in revealing particle acceleration mechanism
in solar flares. A loop-top hard X-ray source was discovered in a compact flare (Masuda et al.
1994), indicating that energy release and particle acceleration occurs above the post-flare loop.
Using a unique technique “time-of-flight method,” Aschwanden et al. (1997) also showed thatthe particle acceleration occurs above the post-flare loop. It is currently believed that particles
are accelerated in the current sheet and slow/fast-mode shocks coupled with the reconnection
dynamics. However, we have not at all understood how particles (ions and electrons) are

1-21

accelerated. This is because of the difficulty in directly probing the acceleration region. One of
the main reasons for the difficulty is the lack of dynamical range in hard X-ray observations.
Accelerated electrons radiate X-ray emissions via non-thermal bremsstrahlung process. The
amplitude and energy of X-ray emissions depends on the number and energy of accelerated
electrons. An impediment to fully understanding of acceleration mechanism is a dynamic range
of instrumentation. The proposed photon-counting grazing-incidence X-ray telescope has
dynamic range three orders of magnitude better than RHESSI. It covers the energy range from
<1 to ~5 keV (goal ~10 keV), which contains thermal emission both in continuum and lines aswell as non-thermal emission at the higher end of the energy range. Suprathermal electrons(~10keV) have most parts of non-thermal energy. Non-thermal signatures have been seen evenbelow ~5 keV in many microflares (Hannah et al. 2008). Therefore, the proposed X-ray
telescope may allow us to evaluate non-thermal electrons.
Another way to discuss the non-thermal energy is by determining whether the electron energy
distributions are Maxwellian or not with the EUV spectroscopic measurements. Seely et al.
(1987) studied the change of line intensity ratios of He- and Li-like Fe ions with
non-Maxwellian electron energy distribution in the laboratory experiment. The bottom twopanels of Figure 1.17 show Fe XIV lines with thermal and non-thermal equilibrium conditions.
They show that intensity ratios marked with red circles changes in these two cases. Differential
emission measure and electron density can affect on the intensity ratios, as well as non-thermal
equilibrium state. Therefore, multiple line observations (see Appendix 3-A) can distinguishwhich affect on the intensity ratio. The Hinode EIS cannot perform such analysis because ofpoor wavelength coverage, but the proposed UV/EUV spectrometer will provide a quantum
leap in our understanding of electron energy distributions and particle acceleration in the solar
atmosphere.

Figure 1.17. Two electron energy distributions corresponding to Te ~ 650 eV (top panel) andchanges of Fe XIV line intensity ratios (bottom panel) highlighted by dotted red circles.

1-221.3.6.

What is the role of 3-dimensionality in magnetic reconnection?

Another important question in solar flares is what is the role of 3-dimensionality in magnetic
reconnection. Generally the current sheet thickness of the magnetic reconnection is in order ofion inertia/gyro radii scale (a few 10m). However, solar flares show a large (a few 100Mm)
well-known flare arcade and ribbons (see Figure 1.19). The difference of scale size between
current sheet thickness and current direction is roughly 7 orders of magnitude. The question isthat the magnetic reconnections take place individually in current direction, or there are somecausal connections for current direction. In other words, we can say that the simultaneous flare
triggers happen in each current sheet and they are developed to a large flare structure, or only
one local flare trigger happens and the former magnetic reconnection causes the latter magnetic
reconnections. To clarify this point, we should investigate that the existence or nonexistence ofthe causal connection for current direction in solar flares.
One of the advantages on using the solar plasma for magnetic reconnection studies is thatimaging observations of reconnection events, such as flares and jets, can provide the global
picture of plasma condition and magnetic field configuration. We already have some indicationfrom modern satellite observation, and many modeling study has been done. According to theirresult, flare arcades must be individually heated as large numbers of distinct loop strands inorder for them to reproduce the observed emission signatures (e.g., Hori et al. 1997; Reeves &Warren 2002; Reeves, Warren, & Forbes 2007; Warren & Doschek 2005). This multi-loop
model provides a good description of a coronal arcade once it has come to rest after a flarereconnection event. Some theoretical studies show that patchy reconnection (Figure 1.18) cancreate a large number of distinct reconnected loops which each descend individually throughthe corona. Although we know that each post-flare loops have different characteristics, we do
not have enough information about the causal connection for each flare loops. To clarify thispoint, we need spectroscopic observation (Doppler shifts and line width) with high temporal
resolution. If there are some relationships between each flaring loops, the information of
triggers may propagate with Alfven speed. Thus at least we need temporal resolution shorter
than Alfven transit time. Therefore, spectroscopic measurements with high cadence (shorter
than Alfven time) can open a new window for investigating the role of 3-dimensionality inmagnetic reconnection.

Figure 1.18. Field connectivity change predicted in patchy reconnection
1-23

1.4. Small-scale Physics Initiating Large-Scale Phenomena
1.4.1.

Flares and CMEs

Solar flares and coronal mass ejections (CMEs) are manifestations of explosive energetic
events at the Sun that accelerate particles. Both sources cause geo-effective particle events that
endanger space assets and human explorers.
1.4.1.1. Coronal Magnetic Fields for Initiating CMEs and Flares
The start of most CMEs is associated with the eruption of a prominence/filament or sigmoid.
Both prominence/filaments and sigmoids are manifestations of non-zero magnetic helicity; thatis, they show evidence of non-potential magnetic fields, and thus electric currents and twisted
magnetic fields. Both have been modeled as magnetic flux ropes (e.g., Rust & Kumar 1996;
Titov & Demoulin 1999). Essentially all models of sigmoid formation involve twisted ropes ofmagnetic flux, since the twist affords a convenient method of storing energy needed for an
eruption, and the appearance of the sigmoids is suggestive of writhing coronal loops(McKenzie & Canfield 2008; Savcheva & van Ballegooijen 2009).
The presence of measurably twisted fields before the eruption facilitates an estimation of thestored energy available for the eruption. The coronal magnetic fields are modeled primarily
with a non-linear force free field (NLFFF) model. The magnetic fields measured at thephotospheric level are currently used for this modeling, but the photospheric data are notconsistent with the NLFF condition. The force-free condition is well justified in low beta
plasma, which is in the corona and the upper chromosphere and is not in the photosphere.
Nevertheless, precise photospheric vector magnetic field data from Hinode is currently themost key observation for NLFFF modeling of coronal fields in data-driven simulation of large
flares and CMEs (e.g., Kataoka et al. 2009). The data-driven simulation will be a powerful
method for understanding what triggers eruptions, flares, and CMEs. It is desirable to establish
a reliable method for measuring the Poynting flux in an active region in order to calculate theavailable free energy for an eruption. Predicting the size of an eruption is important for space
weather applications.
Magnetic fields measured in the chromosphere will be a key for improving the accuracy of
NLFFF modeling. Direct chromospheric magnetic field information, e.g. H images can beused as additional constraint (Wiegelmann et al. 2008). H images show the horizontal
direction of the chromospheric magnetic field vector. Direct measurements of line-of-sight
chromospheric magnetic field help to serve as an additional constraint. Measurements of themagnetic field vector in the low beta chromosphere could be used directly as boundary
condition for NLFFF modeling.
An important issue for NLFFF-modeling is to have sufficient large field of view that
accommodates most of the connectivity within a region and its surroundings (DeRosa et al.
2009). Because of narrow field of view (200 arcsec x 200 arcsec), the high spatial resolution
magnetic field data from Solar-C needs to be imbedded into a large field of view data fromother spacecrafts, such as Solar Dynamic Observatory (SDO), and ground-based observatories.
Therefore, for accurately estimating the magnetic field of flares and CMEs in the corona, it ismost important to coordinate the Solar-C observations with instruments with wide field of view,
such as SDO and ground-based observations.
1-24

1.4.1.2. Small Scale Triggers and Dynamics in Initiation
In order to accurately predict when, where, and how space weather events are initiated, wemust understand the interrelationship between small-scale processes on the solar surface andlarge-scale magnetic structures in the solar corona. In particular, current belief is that
small-scale processes (e.g., magnetic reconnection) are important as a trigger for solar flares.
The solar magnetic fields control the flow of energy and formation of structures above thephotosphere. Hinode observations would make much more progress in understanding the roleof magnetic topology (such as quadrupolar fields in the breakout model) and small-scale
photospheric flux emergence in flare and CME initiation, if Hinode captures many more flare
events in coming years. Flux emergence is one of key magnetic activities in the photosphere toquickly develop magnetic topology for flare and CME initiation.

Figure 1.19. (top panel) Flare ribbons in a Ca II H
image taken by Hinode SOT, with blue contours
showing micro-wave signals recorded with
the Nobeyama Heliograph and green contours
showing hard X-ray signals recorded with
the RHESSI satellite. (right panel) An X-ray image
taken with Hinode XRT before the flare begins.

Measurements of magnetic fields in the higher layer, i.e., chromosphere, and mass flows overthe entire atmosphere will observationally provide new insights in understanding flare and
CME initiation. For example, many flares are related to newly emerging flux, and the higher thecomplexity and shear of the field the greater the chance of flares occurring. We do know that
there is a ‘time lag’ from when flux emerges to when there is a response in the corona (e.g.
Schmieder et al., 2004). Flux emergence can occur 6 hrs before a flare occurs. In the 13th
December 2006 X-flare (Figure 1.19), the increase in coronal turbulence occurred 12 hours
before the flare began (Harra et al., 2009) and 16 hours after the helicity injection rate reached apeak (Magara and Tsuneta 2008). Further work studying the detailed response of the
1-25

chromosphere and corona to emerging flux has shown that the chromospheric response occurs
first (Li et al., 2007). There is also evidence that the corona shows a small increase in flows
before flaring begins (Harra et al., 2010). The chromospheric emission shows smallbrightenings over the complex region of emerging flux. The corona behaves in a completely
different way, showing the development of large-scale loops with enhanced outflows at theedges. The linkage between the chromosphere and corona nor the energy transport processes
that follow the flux emergence are not understood yet. Theoretical simulations have been
carried out to explain flux emergence (e.g. Cheung et al., 2008), but the full connection frombelow the photosphere to the corona and the linkage with magnetic reconnections in corona andchromosphere is not yet achievable.
Flare initiation studies will be partially performed with Hinode, if Hinode captures many
more flare events in coming years. With this in mind, we define the Plan B science goals more
specifically. Hinode can provide accurate measurements of vector magnetic fields at thephotospheric level, but such maps covering a large area are available only with extremely longinterval (several hours or longer). Considering the timescale of emergence of each magnetic
flux (5-15 minutes to cross from the photosphere to the corona), much increase of cadence willallow us to capture details of the flux evolution. Enhanced telemetry resource about one orderof magnitude larger than Hinode will increase mapping cadence to less than several tens ofminutes even for a large field of view (200”x200”) with 0.2” resolution. In addition,measurements of chromospheric magnetic fields can provide the dynamical evolution of
emerging flux and nearby pre-existing flux in 3D domain.
There are some signatures of changes in the corona before a CME erupts, e.g., increasingcoronal flows (Baker at al., 2009). This indicates that the overlying magnetic fields arechanging which reduces the stabilizing forces of the overlying magnetic field. The energy
necessary to produce these changes clearly comes from magnetic fields that emerge from below
the surface, but the mechanism for transporting this energy to regions above the surface isunclear, including its effect on the moment the CME is triggered. Current indications are thatreconnection must be occurring in the lower atmosphere to build up the pre-eruption magnetic
structure of a flux rope (Green and Kliem, 2009). But this cannot be presently modeled due tolack of spectral information in the low atmosphere. EUV/FUV spectroscopy covering from thelow atmosphere to the corona can reveal the entire picture of mass flows.
1.4.1.3. Fine Magnetic Field Structures in Prominences/Filaments
A typical prominence – filament on the disk – is characterized by chromospheric coolermaterial (about 10,000 K) suspended at coronal heights, a configuration that can only persist if
the magnetic field supports the plasma against the force of gravity. Since major flares and
CMEs are sometimes accompanied by an eruption of the prominence, the formation anddynamics of the prominence is also essential to address questions of magnetic energy
storage/release for a CME-like explosion.
Hinode SOT observations have revealed that prominences have very complicatedgeometrical configurations, and they look different in active region and quiet Sun (Figure 1.20);
Active-region prominence consists of nearly horizontal threads (Okamoto et al. 2007, Merenda
et al. 2007), whereas quiescent prominence has apparent vertical threads (Berger et al. 2008),
although magnetic field is supposed to be oriented horizontally. Prominences are also supposed
to have helical magnetic configuration to support cool materials in the corona, but there is no
direct magnetic field measurements to confirm the helical configuration in prominence.

1-26

Direct measurements of magnetic fields in prominences have been required to reveal theoverall magnetic field configuration in prominence and how the configuration is developed.
Weak magnetic fields in prominences produce polarization signals by so-called Hanle effectand therefore magnetic field orientation can be inferred with sub-arcsec spatial resolution when
the polarization profiles of chromospheric lines are measured with the proposed 1.5m telescope(SUVIT). It would be a great progress in understanding the triggers of CMEs and flares, if themagnetic fields can be measured during the eruption of prominences.

Figure 1.20. Prominences observed in an active region (left) and a quiet region (right) with
Hinode. It is important to clarify how the fine structures smaller than 1 arcsec are associatedwith magnetic field configuration.
1.4.1.4. Solar Flare Forecasting
One of the most important goals of solar physics research is to improve our ability to predict
geo-effective solar flares and CMEs. The Solar-C mission has observational capabilities toadvance our understanding and prediction capability for space weather, primarily by exploring
how small-scale physical processes are related to the initiation of space-weather phenomena.
It is empirically known that the occurrence rate of big solar flares correlates with someextensive quantities representing the properties of magnetic field and flow in flaring active
regions. For instance, the probability for big flares to occur within 24 hours exceeds 50% in big
active regions, where the extensive magnetic quantity, e.g. unsigned magnetic flux, is larger
than some critical value. However, the skill score for forecasting flares is still low, and it is too
difficult to predict the onset of solar flares with high reliability (Leka and Barnes 2007, Barnes
and Leka 2008). It is due to the fact that the major fraction of flares occurred in moderate active
regions, where the extensive quantity is lower than the critical value. The result indicates thatthe extensive quantities, which can be obtained by the integration over whole active region,
cannot provide us sufficient information to forecast solar flares, and also suggests that theoccurrence as well as the scale of solar flares is susceptible to the trigger process, which mightproceed in much smaller region than the size of active region. Therefore, to identify the criticalindications of flux emergence that causes flares and CMEs is crucial for developing flareprediction algorithms. If we can discriminate the small-scale activity triggering solar flares interms of the unprecedentedly high resolution observation, it would be likely that thepredictability of solar flares could be improved and that the Solar-C would greatly contributethe space weather forecast.

1-27

1.4.2.

Solar Winds

1.4.2.1. Origins of Fast Solar Winds
Coronal holes (CH) are the source of the fast solar wind (Krieger et al. 1973). For polar
coronal holes, Hassler et al. (1999) found regions of strong outflow lying at the boundary ofmagnetic network cells from observations in the λ770A Ne VIII resonance line formed at thebase of the corona (T=0.63 MK). Outflows are also seen in the much cooler λ584A (Wilhelm etal., 2000) and λ10830A (Dupree et al., 1996) He I lines. Tu et al. (2005) find that the wind starts
flowing between 5 Mm and 20 Mm from the measurement of C IV and Ne VIII, respectively.
However, there is still a substantial lack of knowledge on which structures within coronal holes,
and what physical mechanisms, are primarily responsible for the wind acceleration. Three
candidates have been discussed as the structures responsible for the fast solar wind: plumes,
interplumes, and jets (Figure 1.21).

Figure 1.21. Some possible sources of the fast solar wind (plume, interplume, jet) and slow
solar wind (coronal hole-steamer boundary, edge of active regions, reconnection between open
field lines and close loops, reconnection at streamer stalk).
Being the most prominent feature within CHs, plumes have been and are considered as astrong candidate for the source regions of the fast wind, but theoretical models give opposite
indications; Del Zanna et al. (1997) and Casalbuoni et al. (1999) show that the outflow speed inplumes can be larger or smaller than in the ambient depending on the temperature and Alfvén
wave flux assumed in the two regions. VUV lines formed around 1 MK allow the study of theplumes roots on the disk where their intensity is enhanced with respect to the background CH
emission. In addition to measuring Doppler shift along line of sight, Doppler dimming features
seen in, for example, O VI 103.2/O VI 1037A lines allows to determine the radial component of
1-28

the plasma outflow velocity from the analysis of off-limb spectra, although some assumption
on the geometry of coronal structures is necessary. From SUMER and UVCS combined data,
Teriaca et al. (2003) derived the outflow vs. altitude profile in interplumes. However, the result
for the plumes is controversial with very different results being present in the literature(Giordano et al. 2000; Gabriel et al. 2003, 2005; Raouafi et al. 2007). Jets are often associatedto plums (Raouafi et al. 2008), which may cause different results, because the integration of
signals recorded by the 36 to 48 hours’ SUMER observations is required to derive reliable
results. High throughput performance improved by one order of magnitude from SUMER willallow us to study time variability of outflows at timescale of a few hours. From EUV
spectroscopy, plumes are known to be denser and cooler than the surrounding interplumeregions (e.g., Wilhelm 2006). Spectral lines are observed to be broader in interplumes (e.g.,
Wilhelm, 2006; Banerjee et al. 2000) hinting at preferential energy deposition in thesestructures. Signatures of slow magnetoacoustic waves have also been found in both plumes(DeForest & Gurman 1998, Ofman et al. 1997; 2000) and interplume regions (Banerjee et al.
2000; 2009), and it is important to understand observational link between waves and velocity
structures (see also section 1.2). Moreover, the spectroscopic analysis of chemical composition
could provide a tool to establish whether plumes or interplume regions contribute significantly
to the fast wind by comparing the composition of the different structures on the Sun with in situ
data.

Figure 1.22. Hinode XRT false-color image of the north polar coronal hole. A typical jet is seenin the center of this image (from Cirtain et al. 2007).
The Hinode XRT has revealed a large number of jets occurring within the polar holes (Cirtain
et al., 2007, Figure 1.22). Given the high velocities attained during these events (up to 800km/s), they may be contributors to the wind. The Hinode SOT has revealed wide distribution ofstrong field patches in the polar regions (Tsuneta et al. 2008, Figure 1.23), by viewing the polar
regions at 7 deg inclination periods with 0.3 arcsec spatial resolution. Flux emergence appears
into open fields (e.g., coronal holes), resulting in jets (e.g., Shimojo and Tsuneta 2009). The
standard model for jets associates the energy source with reconnection in the corona, but
detailed chromospheric and transition region diagnostics are currently not available. Current
observations also lack the spatial resolution to see the sub-structure of jets. Jets are important
targets for close examination because they are the most likely place to observe reconnection
and its related sub-structure in the corona. Because of the relatively small scale of jets, areasonably accurate accounting of the magnetic flux budget, and the energy budget, should be
1-29

achievable with high-resolution, high-sensitivity photospheric and chromosphericmagnetograms, paired with high-resolution, high-sensitivity, high-cadence coronal imagery.
Such studies should enable measurements of the energy storage prior to jet eruption, the rate ofenergy release, and the rates of helicity transfer. In light of the Hinode observations of coronal
jets, it is believed that the jets may contribute a significant amount of mass into the fast solar
wind.

Figure 1.23. Strong patches of kG magnetic fields distribution in the polar region, reveraledwith Hinode SOT magnetic field measurements (Tsuneta et al. 2008).
1.4.2.2. Origins of Slow Solar Winds
The source of the slow solar wind is even more difficult to identify than that of the fast solar
wind. Sources include coronal holes boundaries, edges of the streamer belt and edges of active
regions (Figure 1.21). The proposed UV/EUV spectrometer is capable of measuring relative
element abundances in different atmospheric structures that can be related to in-situ
measurements from Solar Orbiter and Solar Probe+. The abundances in the slow solar wind are
different from in the fast solar wind. Thus, abundance measurements can be used with in-situ
measurements to locate the origins of the slow solar wind. The Hinode observations haverevealed the persistent flows at the edges of active regions (Figures 1.24, Sakao et al. 2007).
These could provide up to 25% of the slow solar wind. Recent STEREO and Wind observations
suggest that a considerable fraction of the slow solar wind is transient in nature even at solar
minimum – these have small scale-sizes and low magnetic field strength. They seem to berelated to coronal source surface sector boundaries (Kilpua et al., 2009) that could be dynamicalchanges at coronal hole boundaries or edges of helmet streamer belts. The process involved is
not fully understood and it is widely debated whether these outflows are caused by waves orreconnection. Certainly the magnetic configurations involved suggest large-scale magnetic
reconnection, however there is also some evidence for wave-like behavior of the outflow. The
in-situ composition data for these transient slow winds seem consistent with open-field regions,1-30

rather than closed-field regions. The proposed UV/EUV spectrometer can measure abundances
throughout the atmosphere and map the differences and determine how the abundances evolve.
The active region upflow speeds derived from spectral line wavelengths vary with spectral
lines formed at different temperatures, which implies that the upflows consist of fine scale
structures with difference speeds. At present it is impossible to connect these flows withtransition region and chromospheric structures, due to the limited wavebands and spatial
resolution of Hinode EIS. If we are going to understand how these flows arise, i.e., the source of
their energy, and where in the atmosphere they originate, a higher spatial resolution instrumentwith connectivity to lower temperatures is essential.
To make further progress, it is necessary to sort out properties of the solar wind source
regions among a variety of possible source structures. It is important to ‘follow’ the field lines
from the chromosphere, through the transition region, up into the corona by comparing andrelating plasma properties such as density, temperature, line width, Doppler shift and elemental
abundance, in space and time across the footpoint region of slow wind. At the same time,magnetic field extrapolations are needed to identify preferable location of magnetic
reconnections in the transition region and corona and to show the connections to theheliosphere.

Figure 1.24. Persistent plasma outflows at the edges of action regions, revealed with Hinodeobservations.

1-31

Chapter 2:

Solar Ultra-violet Visible and IR Telescope

2.1. Introduction
2.1.1. Solar UV-Visible-IR telescope in SOLAR-C Plan-B
The emergence of magnetic fields, its forcing by convective flows on the solar
surface, and its interaction with previously-emerged magnetic flux, give rise to energetic
phenomena in the solar atmosphere through magneto-hydrodynamic processes. The
Hinode Solar Optical Telescope (SOT) now provides the first high-resolution
diagnostics of the solar magnetic field from a space platform, and now allows
discovering the processes by which magnetic flux interacts with the convective flows in
the lower solar atmosphere, “photosphere”. Furthermore, Hinode and recent
ground-based telescopes has provided strong indications that those photospheric
dynamics initiate outward propagating waves and dynamic reconnection events that
result in acceleration and heating of the plasma accompanied by significant outward
energy and mass flux in the upper atmosphere, i.e. “chromosphere” and “corona”.
Unfortunately, our knowledge on the magnetism of the solar chromosphere is practicallyinexistent notwithstanding the qualitative information provided by monochromatic
images at various spectral lines sensitive to the chromosphere. Such intensity images do
not provide any quantitative information on the temperatures, velocities, and magnetic
fields.
In these circumstances, the primary objective of the instrument, Solar Ultra-violet
Visible and IR telescope (SUVIT), is to explore dynamical nature of the magnetic solar
atmosphere and physical processes responsible for making the dynamics throughspectroscopic and polarimetric measurements of physical conditions (e.g. temperaturesand velocities) with emphasis on magnetic fields in the chromosphere. Quantitative
information on magnetic fields in the place where magnetic energies are transferred andliberated is critically important to expand our knowledge on the magnetized atmosphere
although it has been significant challenges in measuring magnetic fields above thephotosphere because of lower polarization signals in spectrum lines emanating from thechromosphere.
The observations also demand parallel efforts in advanced data analysis,
interpretation, and modeling methods. Fortunately, at present we live in a time whererapid advancement in the analysis and modeling are taking place, and the models
describe their synthetic chromosphere/corona with increasing realism (see Appendix
A.3 and A.4). These theoretical advances demand new observational constraints. The
present concept of SUVIT is structured to provide those specific observational
constraints.

2-1

2.1.2. Science targets with the instrument


Magnetized solar atmospheres: how they are created and evolve.

Hinode has demonstrated that aspace-based telescope has strong
benefits
inprecisespectropolarimetric observations. It allows
studying structures and dynamics inthe photosphere, such as formationof strong flux concentration, its
interaction with granular convection,andsunspotsstructures
anddynamics associated with strong
magneto-convection. The next stepis to extend our quantitativeknowledge of the solar atmospherefrom the photosphere to the Fig. 2.1. SUVIT on SOLAR-C extends our quantitative
chromosphere (Fig. 2.1). In the knowledge on the magnetized solar atmosphere from thechromosphere, magnetic pressure photosphere to the chromosphere.
becomes comparable with or
dominant over gas pressure (i.e.
plasma beta ≤ 1) while gas pressure is dominating in the photosphere. This is why
measurements of magnetic fields in the atmospheric layer higher than the photosphere isdemanded, which also helps to infer magnetic field configuration in coronae throughmagnetic field extrapolation such as potential field and force-free field modeling.
In the quiet Sun, granular-scale horizontal field emergences are ubiquitously found in
the photosphere, which were extensively studied using Hinode (Lites et al. 2008,
Centeno et al. 2007, Orozco Suarez et al. 2007, Ishikawa et al. 2008). An outstanding
question is how much fraction of magnetic fields keep expanding into the upper
atmospheres, and how they are reconfigured to create canopy structures overlying thequiet photosphere. The observations are truly essential to know eventual history of
emergent magnetic fields. High-resolution observations with Hinode and recent
ground-based telescope showed super-sonic upflows and downflows ubiquitously seenat the photospheric layer. Some of them are associated with formation of strong flux
tubes (e.g. Nagata et al. 2008), and some of them are associated with cancellation ordisappearance of bipolar magnetic flux (e.g. Kubo et al. 2010). Driving mechanisms andultimate consequence of the flows could be physically interpreted with SUVIT owing to
combined measurements of plasma properties and magnetic fields at the multiple
atmospheric layers.
A solar prominence is also an important target with SUVIT because it carries
significant amount of magnetic flux and masses in coronal mass ejections (CME). In
spite of the importance, its magnetic field configuration is poorly known. This is
primarily because of lack of observational information on magnetic field configurationin the prominence. Hinode SOT observations revealed that structures in prominences
are very complicated in the spatial scale smaller than 1000 km (Okamoto et al. 2007,
2-2

Berger et al. 2008). It is essential to understand how the fine scale structures are createdand how cool materials are provided and sustained in the upper atmosphere. There aresome observations of field emergence at the photospheric layer with Hinode (Okamoto
et al. 2009, Lites et al. 2010), but long-term observations of the photospheric andchromospheric fields with SUVIT provide an entirely new view on formation andevolution processes of the prominences.


Small-scale dynamic phenomena responsible for the energy transfer.

One of the science goals of the Solar-C plan-B is to elucidate the energy and masstransfer from the photosphere upward through the chromosphere to the corona.
Although the chromosphere occupies only a very thin region in physical dimensions
relative to the structure of the overlying corona, it is the region where most of the
non-radiative heating takes place, and it forms the conduit for the upward flux of massand energy.
Significant progress has already been realized in understanding the chromosphere(Fig. 2.2). Above the very quiet photospheric internetwork where the magnetic field is
weakest, the chromosphere is dominated by largely acoustic shocks initiated in thephotosphere and below by action of the solar convection and the p-mode oscillations(Carlsson and Stein 1997, Rutten 2007). Recent observations from Hinode and other
instruments have suggested that the network fields, interacting with convection,energize and structure the chromosphere. In particular, the “type II spicules” now appear
to be a candidate for channeling the upward movement of mass and energy toward thecorona (De Pontieu et al. 2011). Active phenomena in the chromosphere increase within
active regions, where outward propagating jets, such as chromospheric micro-jets inpenumbrae (Katsukawa et al. 2007), Ellerman bombs (Georgoulis et al. 2002) andchromospheric jets (Shibata et al. 2007) around sunspots, occur frequently. In most of
the active phenomena in the chromosphere, magnetic reconnection is suggested to play
an essential role to energize the outward flow. But there have been no direct
measurements of magnetic field configuration around the reconnection site.
High-resolution observations from SUVIT would provide observations of thechromospheric magnetic field in and around these active chromospheric phenomena,
and also provide the accompanying thermodynamic conditions therein. It is demanded
to get knowledge on what the condition is to trigger the magnetic reconnection and how
fast it dissipates magnetic energies.
Energy transfer between the atmospheric layers takes place through not only radiationand flows, but also propagation of magneto-hydrodynamic (MHD) waves. Though it
has been difficult to catch signatures of waves in observational data, transverse
oscillation of chromospheric structures are now ubiquitously discovered in imaging
observations with Hinode thanks to its high and stable image quality (Okamoto et al.
2007, De Pontieu et al. 2007). In the photosphere signatures of propagation andreflection of MHD waves are found in spectro-polarimetric observations (Fujimura andTsuneta 2009). Quantitative measurements of fluctuations in magnetic fields andthermodynamic parameters are essential to understand how and where the waves are
generated, reflected, and dissipated. Especially it is important to know nature of waves
by identifying MHD wave modes and mode conversion for subsequent shock formation,
resulting in plasma heating.
2-3

Fig. 2.2. Structures and dynamics seen in the photosphere and the chromosphere. In thechromosphere, lots of dynamics are taking place in the broad spatial and temporal range. Velocities
of waves and flows are typically 10 – 100 km/s, one order of magnitude faster than typical velocities
seen in the photosphere.

2.2. Basic requirements to the instrument
The science targets described above leads to basic requirements to the instrument assummarized below:

Polarimetric sensitivity and accuracy

The polarimetric sensitivity in Hinode/SOT and existing ground–based telescope istypically ~10-3, which is enough to get polarization signals with photospheric lines but
marginal to detect the polarization signals in the chromospheric lines. For detailed
measurements of the circular polarized profiles induced by the longitudinal Zeeman
effect and linear polarized profiles induced by the Zeeman and Hanle effect in the fieldstrength down to 10 gauss, we need the sensitivity of at least 10-4 (see A.2 and Fig. 2.23),
which means we have to collect at least 108 photons when the sensitivity is limited withphoton noise. In order to collect the number of photons with 0.2” spatial sampling and
10 to 20 sec temporal sampling, we need an aperture diameter larger than 1 m (see Fig.
2.3). It is necessary to have a strategy to minimize instrumental polarization and to
remove residual polarization cross-talk with pre-flight and in-flight calibration toachieve high polarimetric accuracy.

2-4

Fig. 2.3. Collectable number of photons as a function of the telescope diameter at major
chromospheric lines (right). It is assumed that the overall efficiency of the instrument is 0.1 with thespatial sampling of 0.1”, the temporal sampling of 1 sec, and the wavelength sampling ofΔλ/λ=2x105. Diffraction-limited spatial resolution as a function of the telescope diameter (left).

Spatial resolution

In order to resolve the structures in the chromospheres such as spicules and fibrils,
whose spatial scale is smaller than 0.2”, and also to detect small amplitudes ofmagneto-hydrodynamic waves and changes of magnetic fields associated with magnetic
reconnection taking place inside the chromospheric structures, we need 0.1” resolution
at the shortest wavelength (Fig. 2.2) and diffraction-limited resolution in the wavelength
range longer than visible lights (0.18” at 1.1 μm in the case of ϕ 1.5 m, Fig. 2.3). In
order to collect plenty number of photons required for the polarimetric sensitivity of
10-4, we might have to mitigate the spatial sampling while keeping spatial sampling of0.2” to 0.3” required to dissolve the elemental chromospheric structures.
Spectrum lines

The He I spectrum line at 10830 Å provides the best magnetic field diagnostic byjoint action of the Hanle and Zeeman effect with less affected by atmosphericstratification in the photosphere and lower chromosphere. The Ca II line at 8542 Å isthe second prioritized one for magnetic field diagnostics because it has good sensitivity
to the Zeeman effect in a wide range of magnetic scenarios. Simultaneous observationsof a companion line, He I D3 at 5876 Å for He I 10830 Å and Ca II at 8498 Å for Ca II
8542 Å, are optional to help robustness of diagnostics. Simultaneous polarimetricmeasurements of a photospheric line are also important for identifying energy sources in
the photosphere. There are photospheric lines available near the chromospheric lines, FeI 8538 Å in the Ca II 8542 Å window and Si I 10827 Å in the He I 10830 Å window.
Mg II h and k lines at around 2800 Å are the best line to capture temperatures andvelocities of fibril structures in the upper chromosphere because of higher opacity than
Ca II lines. Because the lines are not accessible in a ground-based observation, they canbe powerful tools for the diagnostics of the high atmosphere with SUVIT. See 2.3 formore details on selection of the spectral lines.
2-5Temporal resolution

We’ve realized ubiquitous occurrence of plasma ejection (i.e. jets) in the chromosphereand MHD waves pervading there. It is critical to resolve transit of jets and waves along
and across the chromospheric structures (Fig. 2.2). The time scale of the dynamic
phenomena regulated by magnetic fields is typically scaled with Alfven transit time, andis shorter than a few seconds when the Alfven velocity is 100 km/s and the spatial scaleof 200 km (corresponding to 0.3 arcsec). In order to collect plenty number of photons
required for the polarimetric sensitivity of 10-4, we may have to mitigate the temporal
sampling while keeping resolution of 20 sec (spatial scale of 200 km divided by thesound speed of 10 km/s) required to dissolve the acoustic travel time across theelemental structures.


Field-of-view

When we use a 4K x 4K format with 0.05” sampling (critical sampling for the 0.1”
resolution) on a detector, a field-of-view (FOV) becomes 200” x 200”, which
compromise most of medium-sized active regions. In force-free extrapolation of coronalmagnetic fields, the FOV provides good coverage to reconstruct complex magnetic
configuration at the core of a flare-productive active region though it is not enough tocover entire magnetic connectivity around the active region.
Simultaneous areal coverage in spectroscopy is also desired to track rapid
chromospheric dynamics. This can be achieved by utilizing a filtergraph observation
simultaneous with the spectrograph requiring slit scanning across the FOV. When the
duration of the plasma ejection is 3 minutes with a velocity of 100 km/s, the plasma
travels 30,000 km corresponding to 40 arcsec, which is the minimum requirement of the
FOV to the narrow-band fieltergraph for imaging-spectroscopic observations.
Spectral resolution

In the detailed spectro-polarimetric observation of the chromospheric lines, theminimum spectral resolution required to resolve a thermal widths of chromosphericlines is λ/Δλ> 100,000 corresponding to 3 km/s in a Doppler shift. In order to observe
spectrum lines emanated from the photosphere, λ/Δλ> 200,000 corresponding to a
Doppler shift of 1.5km/s is required.

2.3. Spectral lines compilation
The main properties of candidate chromospheric lines to be observed with SUVIT canbe found in Table 2.1. The compilation includes spectral lines in the near-UV, visible,and near-infrared wavelength range. Several criteria have been taken into account tomake the line list: they have to be formed within the chromosphere to allow thediagnosis of low β plasma; the lines should show enough sensitivity to plasma
temperatures, velocities, and magnetic field vectors; and the line profiles should beinterpreted through radiative transfer calculations.
2-6

Table 2.1. Candidates of chromospheric lines for SOLAR-C.
Spectral line

1.2λ/D[arcsec]

Photon budget[ /Dl/0.1"☐/sec]
ContinuumCenter4

5

Response height
[km]
Te
V

Zeeman effectHanle
effectgeff

Blong [G]Btrans [G]

Mg II k 2796 Å
Mg II h 2803 Å

0.046
0.046

2.5x102.8x1042.6x102.2x105

~1300
~1200

~1400
~1300

1.17
1.33

85
73

2100
1800

X
-

Ca II K 3933 Å
Ca II H 3968 Å

0.065
0.065

5.5x1056.4x1052.9x105
3.0x105

~1000
~900

~1200
~1100

1.17
1.33

62
51

1500
1200

X
-

H I β 4861 Å
Mg I b3 5167 Å

0.080
0.085

1.1x107
1.3x1071.8x106
1.3x106

~1500
~500

~1000
~600

1.00
2.00

16
3.7

1200
220

X
X

Mg I b2 5172 Å
Mg I b1 5183 Å

0.085
0.086

1.3x107
1.3x1071.0x106
9.2x105

~500
~500

~600
~600

1.75
1.25

6.5
9.6

250
360

X
X

Na I D2 5890 Å
Na I D1 5896 Å

0.097
0.097

1.5x1071.5x1077.2x1058.4x105

~400
~400

~600
~500

1.17
1.33

4.7
3.4

290
230

X
-

H I α 6563 Å
Ca II 8498 Å

0.108
0.140

1.5x1071.1x1072.8x106
3.3x106

~1500
~700

~1300
~1000

1.00
1.07

9.7
6.3

680
300

X
X

Ca II 8542 Å
Ca II 8662 Å

0.141
0.143

1.1x1071.1x1072.1x106
2.1x106

~800
~700

~1200
~1100

1.10
0.83

6.7
8.4

340
430

X
X

He I 10830 Å (QS)He I 10830 Å (AR)

0.179
0.179

1.8x1071.8x1071.7x1071.5x107

~1800
~1800

~1800
~1800

1.42
1.42

18
9.3

380
270

X
X

The list contains the following information (from left to right):
1) Element and central wavelength (λ) of the transition. QS and AR stand for quiet sun and active region
scenarios for the He I 10830 Å line.
2) Spatial resolution for each spectral line at the diffraction limit of a telescope of D=1.5m.
3) Photon budget estimation based upon the preliminary design of Solar-C instrument:- The telescope consists of 5 mirrors coated with Al+MgF2.
- Spectrograph efficiency of 0.2.
- Detector efficiency of 0.5 at λ<300nm, 0.8 at λ>900nm, and typical CCD 300nm<λ<900nm.
- Spectral resolution of λ/Δλ=2x105 (corresponding to 1.5km/s), with Δλ the wavelength sampling.
- Spatial sampling of 0.1" per pixel.
- Temporal sampling of 1sec.
4) Geometrical height at which the sensitivity (response) of the spectral line core to temperature andvelocity perturbations is dominant. The heights were calculated using the centroid of the corresponding
response functions in the FALC atmospheric model (Uitenbroek 2006).
5) Sensitivity limit (1σ) to magnetic fields through the Zeeman effect. The quantities Blong and Btransstand for the longitudinal and transverse component of the field vector detectable with the sampling
shown above (/Dl/0.1"☐/sec ) and geff is the effective Landé factor.
6) The check boxes indicate that the spectral line is sensitive to magnetic fields through the Hanle effect.

2-7

Fig. 2.4. Expected photon budged to be achieved with the 1.5mϕ telescope at major chromosphericlines. See Table 2.1 for details.

Fig. 2.5. Sensitivity of the spectral lines listed in Table 2.1 to magnetic fields through the Zeeman
effect. The lower field strength means higher sensitivity.

2-8

2.3.1. Spectral lines for diagnostics of chromospheric magnetic fields
The sensitivity limit for magnetic field diagnosis through the Zeeman effect can befound in Table 2.1 and Fig. 2.5. The Na I D and Mg I b lines show good sensitivity tomagnetic fields through the Zeeman effect. However, these two lines are formed low inthe solar atmosphere (400 - 600 km). Therefore, they are not appropriate to diagnose themagnetic properties of the true chromospheric plasma. By contrast, the Mg II k/h and
Ca II K/H lines form high in the chromosphere but their sensitivity to magnetic fields
through the Zeeman effect is limited because of the short wavelengths.
The He I 10830 Å line has good sensitivity to the Zeeman effect in both, active andquiet Sun regions, although the sensitivity tends to diminish in the quiet Sun because theresidual intensity decreases, i.e., the spectral line shape becomes shallow. In addition,
the interpretation of the He I 10830 Å triplet is relatively easy because it is formed
within a thin slab in the atmosphere. Hence, the line formation is less affected by thethermal stratification of the photosphere and lower chromosphere. Moreover, the He I
10830 Å triplet can be used for Hanle effect diagnostics. Recently, theoretical tools tomodel the joint action of the Zeeman and Hanle effect in the He I 10830 Å triplet havebeen developed satisfactorily, such that it is very powerful for magnetic field
diagnostics of chromospheric and coronal structures (e.g., forward modeling andinversion code HAZEL developed by Asensio Ramos, Trujillo Bueno & Landi
Degl’Innocenti 2008, see Appendix A.3.2). The longitudinal Zeeman effect at the He I
10830 Å line has a sensitivity down to 10 to 20 gauss with 0.1” spatial sampling andone second integration. We can measure transverse fields down to 0.1-100 gaussthrough the Hanle effect when we perform deep integration (e.g. 0.2” spatial samplingand 20 sec integration).
Table 2.2. Sensitivities (1σ) to magnetic fields at He I 10830 Å and Ca II 8542 Å.

He I 10830 Å(QS)
He I 10830 Å(AR)
Ca II 8542 Å

Blong

0.1 arcsec sampling
1 sec integration~20 gauss

Btrans~380 gauss (Zeeman)Blong~10 gauss

Btrans~270 gauss (Zeeman)

Blong
Btrans~7 gauss~340 gauss (Zeeman)

0.2 arcsec sampling
20 sec integration~2 gauss~130 gauss (Zeeman)
0.1 – 100 gauss (Hanle)~1 gauss~90 gauss (Zeeman)
0.1 – 100 gauss (Hanle)~0.8 gauss~ 110 gauss (Zeeman)(*) Here we assume the same instrument throughput and wavelength sampling used in Table 2.1.

The Ca II infrared triplet at about 850 nm is the second prioritized line for magnetic
field diagnostics in the SUVIT on SOLAR-C. It shows good magnetic sensitivity in awide range of magnetic scenarios, including quiet regions, which potentially remedies
the weakness of the He I 10830 Å line. The longitudinal Zeeman effect especially has agood sensitivity at the Ca II 8542 Å line, and we can measure magnetic fields weaker
than 10 gauss with 0.1” spatial sampling and one second integration. Of especial interestis that through the Hanle effect the linear polarization of this line is sensitive to theorientation of the magnetic field in the upper chromosphere of the quiet Sun (Manso
2-9

Sainz and Trujillo Bueno 2010) though the polarization amplitude due to the Hanle
effect is relatively small (order of 10-4).
In Table 2.2 are summarized the sensitivities of the spectrum lines above mentioned
to magnetic fields through the Zeeman and Hanle effects.
2.3.2. Spectral lines for diagnostics of chromospheric dynamics
The chromosphere is pervaded by fast moving plasma structures. The measure of the
dynamic properties of such structures requires not only spectral lines with great
sensitivity to temperature variations at certain wavelengths but also short integration
times. Since such observations are done typically with filter instruments with limited
spectral bandwidths little contamination from nearby wavelengths is also desired. In thissense, the Mg II k/h spectral lines at around 2800 Å are suitable for the SUVIT in
SOLAR-C. Their sensitivity to temperature variations is great in the upper chrmospherebecause the opacity of these lines is very large at the line core. High temperaturesensitivity lines also allow investigating heating processes in the chromosphere.
Although the low photon budget at the spectral range of the Mg II lines is a limiting
factor for the polarimetric measurement (Table 2.1 and Fig. 2.4), it is well enough to dohigh cadence spectroscopic observation (without polarimetry) with exposure times
shorter than 0.1 sec.
Other spectral lines like the Ca II infrared triplet at about 850 nm and Ca II H and K
lines are good candidates for tracking the dynamics of the various chromospheric events,
by analyzing observed line profiles along with cloud modeling and/or forward modelingof the lines obtained from numerical simulation of the chromosphere (see Appendix
A.3.1 and A.4).
2.4. Configuration of the instrument
A basic concept of the optical design of SUVIT on SOLAR-C is to use a similar
configuration with the Solar Optical Telescope (SOT) on HINODE, which allows us tomake use of the technical heritage that the Japanese solar physics group developed in
HINODE/SOT (Tsuneta et al. 2008, Suematsu et al. 2008). A major difference is toexpand the diameter of the aperture from 0.5 m to 1.5 m, which is aiming to collect oneorder of magnitude more photons than HINODE/SOT as well as angular resolution
better than 0.1 arcsec. This is essential to achieve the high polarimetric sensitivity forchromospheric diagnostics. SUVIT consists of the telescope and focal plane instruments.
The focal plane instruments are installed into two packages:
1)

Spectrograph package
It consists of an Echelle grating spectrograph for spectro-polarimetric observations
and a correlation tracker for the image stabilization system.

2)

Imager package
It consists of a broadband filtergraph utilizing interference filters formonochromatic imaging observations and a narrowband filtergraph utilizing atunable narrowband filter for imaging spectrograph allowing to make
Dopplergrams and magnetograms.

2-10

The combination of the multiple instruments enables to cover wide variety of scientifictargets. The optical block diagram of the instruments is shown in Fig. 2.6. The telescopecollects incident lights and feeds into the focal plane packages. The tip-tilt mirror in thetelescope is driven to remove attitude jitter of the spacecraft using signals of image
displacements detected with the correlation tracker in the spectrograph package.
Polarization modulators are to be installed in each focal plane package whilepolarization calibrator is possibly implemented in the telescope. Preliminary parametersof pixel scales and typical exposures are summarized in Table 2.3. Table 2.4 shows
possible choices of the spectrum lines at each focal plane instrument. The combinationof the spectrum lines provides the spectrum diagnostics completely covering from thephotosphere through the middle to the upper chromosphere.

Fig. 2.6. Optical block diagram of the focal plane instruments.

Table 2.3. Configuration of the focal plane instruments.
FOV
Pix scaleExp.
Note
Correlation tracker (CT)10”x10”0.2”< 2x10-3 sec faster readout than 5kHz.
50x50 format
Spectro-polarimeter184”×184”
0.09”~1 sec2 pix sampling of 0.18”res(SP)
2K pix along slit, S/N~103.
0.18”>10 sec 2x2 pix binning to reach
S/N~104.
Broadband filtergraph164”×164”0.04”<0.1 sec 2.5 pix sampling of 0.1” res(BF)
4K×4K detector(1)
Narrowband filtergraph205”×205”
0.05”<1 sec2 pix sampling of 0.1” res(NF)
120”x120”(2)4K×4K or 2Kx2K detector(1) The optical configuration allows us to cover the FOV of 200” x 200” if a tunable filter has an aperture big enoughto cover the FOV.
(2) The 120” x 120” FOV is available when the diameter of a Lyot filter-type tunable filter has a diamete of 50 mm at
the F/40 telecentric focus. In order to accommodate larger FOV coverage in the NF, we need optics to convert
magnification in the channel.

2-11Table 2.4: Spectrum lines to be observed in SUVIT on SOLAR-C.
Spectrum lines
Mg II k279 nmContinuum
TBD
CN band
388 nmCa II K393 nm
Mg I b517 nmNa I D589 nm
525 nm
Fe I
or 630 nm
HIα656 nmCa II
854 nmHe I
1083 nm

SP
x

BF
x
x
x
x

x
x

NF

x
x

PurposeT and V in the chromosphere
Imaging and T in the photosphere
Magnetic elements in the photosphere
Structures in the chromosphere
V and B in the low chromosphere
V and B in the low chromospherex

V and B in the photospherexxStructures in the chromosphere
T, V, and B in the chromosphere
V and B in the chromosphere2.5. Optical telescope assembly
2.5.1. Optical design
The optical telescope assembly (OTA) of SUVIT employs the following basic
characteristics in the baseline design for exploring its technical feasibilities:- Aplanatic Gregorian 2 mirror telescope.
- Entrance aperture with the 1.5mφ diameter.
- Heat rejection mirror at the primary focus.
- Axial symmetric optics to give negligible instrumental polarization.
- Afocal interface between the telescope and the focal plane instruments.
The aperture size of the telescope is the most critical item in the concept of the telescope,
which is driven by the scientific requirements. We are going to optimize the aperturesize down to 1m according to future investigation of technical and scientific feasibilities.
Another factor to design the telescope is a limitation of the length of the telescope to fit
within a fairing of the launcher (JAXA H-IIA rocket). Based on a study of positional
tolerance of the secondary mirror, the distance between the primary mirror (M1) and thesecondary (M2) is decided to be 2.8 m as a baseline design. Basic parameters of the
optical design of the telescope are summarized in Table 2.5, and its overall
configuration is shown in Fig. 2.7. The afocal interface has an advantage to relax thepositional tolerance between the telescope and the focal plane instruments. Because thewider wavelength coverage (280 nm to 1.1 μm) than Hinode is required in SUVIT-OTA,
the refractive optical design to make a collimated beam in HINODE/SOT is not
appropriate. Instead, we are going to realize it using a collimator mirror unit (CMU),
which consists of three off-axis aspheric mirrors. The third mirror of the CMU can beused as a tip-tilt mirror for the image stabilization system. Optical performance of thetelescope is shown in Fig. 2.8 as is evaluated with spot diagrams at the image plane
focused with an ideal lens located at the exit pupil. The RMS spot radius is well within
the spot size due to diffraction within the field of 200 arcsec.
There are two field stops in the telescope. One is a heat dump mirror (HDM) at the
2-12focus of M1, which is 45-deg flat mirror with a central hole passing the beam of field ofview of 300 arcsec (TBD), and to reflect the unnecessary solar illumination outside of
the hole from the side of the telescope. The other is a secondary field stop (2FS) at theGregorian focus, which is again to reflect the unnecessary solar illumination out of the
field of 220x220 arcsec (TBD), which a bit oversizes the area of detectors in the focal
plane instruments.
Table 2.5. Basic optical parameters of SUVIT-OTA.

Entrance pupil (mm)M1 – M2 distance (mm)M1outer diameter (mm)clear aperture (mm)curv. radius (focal) (mm)conic const.
M2
outer diameter (mm)clear aperture (mm)curv. radius (focal) (mm)conic const. (mm)HDM outer radius (mm)Eff. f-length at 2nd focus (mm)
CMU focal length (mm)Exit pupil diameter (mm)(note: The unit is in mm.)

SOT-OTA(HINODE)
500
1500
560
509
2339.4 (1169.7)-0.9706
159
147
524.94 (262.47)-0.3996
32.83
4527 (F/9.05)
271
30

SUVIT-OTA(SOLAR-C)
1500
2800
1580
1513
4629.34 (2314.67)-0.990927
340
320.5
835.494 (417.747)-0.54815
43.1
14300 (F/9.54)
572
60

Fig. 2.7. The optical configuration of the telescope assembly of SUVIT. A polarization calibration
unit (PCU) can be inserted after the Gregorian focus for calibrating polarization responses of the
focal plane instruments.

2-13

Fig. 2.8. Spot diagrams of the telescope evaluated at the image plane focused with the ideal lens
located at the exit pupil. Note that the RMS spot radius due to diffraction is about 6 μm at 500 nm.

2.5.2. Mirror coating and photon budgets
In order to observe the spectrum lines prioritized based on the scientific point of view,
it is essential to cover the wavelength range from 280 nm to 1.1μm. A coating design ofthe mirrors is critical not only to have higher throughputs for precise polarimetricmeasurements but also to minimize absorption of the solar illumination at the mirrors of
the telescope. The silver-based reflective coating used in HINODE/SOT has highreflectivity in the visible and NIR wavelength range. But its reflectivity is very low in
the UV wavelengths shorter than 350 nm. Aluminum coating is good alternative in UV
although its reflectivity is slightly lower in visible and NIR (Fig. 2.9).
When we apply the aluminum coating to the mirrors in SUVIT-OTA, and if we
assume efficiencies of the focal plane instruments, we can estimate the number ofphotons to be collected by the telescope per a resolution element and per unit time. Fig.
2.10 shows effective areas and efficiencies as a function of the wavelength, in which we
assume that the spectrograph or filtergraph efficiency is 0.2, the quantum efficiency ofthe detector is 0.8 at λ>900nm, 0.5 at λ<300nm, and typical QE of CCD at
300nm<λ<900nm. The instrument has a total efficiency of about 0.1 in the observing
wavelength range even when we use the aluminum coating, and the effective area is
expected to be 1000 to 2500 cm2 dependent on the wavelength. Fig. 2.11 shows thenumber of photons to be achieved with the telescope at major chromospheric lines
2-14

where we assume wavelength sampling of λ/Δλ=2x105, spatial sampling of 0.1” /pixel,
and temporal sampling of 1 sec. The polarimetric diagnostics of chromospheric lines
require the polarimetric sensitivity of 10-4. This means that we have to collect 108
photons in a resolution element if the sensitivity is limited by a photon noise. The figure
indicates that we can collect more than 106 photons at the line center of each spectrum
line with the sampling mentioned above. It appears that we can collect 108 photons at
the line center of He I 10830 Å and Ca II 8542 Å if we apply 2x2 pixel summing (i.e.
0.2'' per pixel) and integrate for 10 to 20 seconds, that is short enough to temporally
resolve the dynamics in the chromosphere.
As is shown above, the 1.5mϕ aperture and the aluminum coating provides a good
photon budget satisfying the scientific requirements at observing spectrum lines. For
further improvements of the throughput and for helping a thermal design by minimizing
the heat absorption at the mirrors of the telescope (see 2.5.4), it is desirable to develop amirror coating with higher reflectivity in visible and NIR while keeping the reflectivity
at 280 nm.

Fig. 2.9. Mirror reflectivities of the silver (Ag) coating and the aluminum (Al) coating with MgF2
over-coating.

Fig. 2.10. Effective area of the SUVIT with the 1.5 mϕ aperture. The telescope consists of 5 mirrors
with Ag coating (red line) and Al + MgF2 coating (blue line).
2-15

Fig. 2.11. The number of photons to be achieved with the telescope at observing chromospheric andphotospheric lines. The gray bars represent the number of photons at the line core of each spectrum
line while the white bars are at nearby continuums.

2.5.3. Structure design
The structural design of the
SUVIT-OTA basically follows that of
Hinode SOT. M1 and M2 for theaplanatic Gregorian are made oflight-weighted low expansion material ofeitherULE
orZerodur.
The
light-weighting technology of telescope
mirrors has been much advanced since
Hinode OTA manufacturing era. We
expect 90 % mass reduction with an open
back mirror structure. Then the M1 of 1.5
m aperture would weigh less than 150 kg.
The light-weighted mirrors are crucial not
only from tight weight budget of satellite
but also from ground testing of telescope
optical performance because of small
gravitational mirror deformation. The
surface figures of the mirrors are
ellipsoids as conic constants shown in
Table 2.5. Since null correctors for these

Ring plate

Center section

Top ring
Upper truss

Lower truss

Mirror cellFig. 2.12. CFRP truss structure of SOT-OTA in
HINODE.

2-16

aspheric mirrors can contribute to inaccurate surface figures, we should include a
M1-M2 combination test in the mirror polishing process. We note that flight-model
supporting mechanism for M1 also should be used in the combination test so that thesurface error due to the mounting can be corrected in the polishing process. The HDM
and the 2FS at the primary and the Gregorian focus are to be made of aluminum alloy
and to be coated with a silver coating.
The short Gregorian telescope demands small static mis-alignment tolerances for M1
and M2, on the order of a few tens microns for decenter and despace or several arcsec
for tilt, and a micron-order de-space short-term stability on-orbit during observations(Table 2.6). To meet this tolerance, the telescope structure is made of a truss ofultra-low-expansion CFRP (Carbon Fiber Reinforced Plastics) pipes in a Graphite
Cyanate matrix, whose CTE is proven to be smaller than 0.1 ppm/K, and thedimensional change due to moisture absorption is about 30 ppm, which is much smaller
than conventional epoxy matrix composite pipes (Fig. 2.12). CFRP honeycomb
sandwich panels are adhesively bonded with upper and lower truss pipes without any
metal junctions to save weight and also to avoid differential CTE, which may cause
unexpected thermal distortion.
Table 2.6. Mis-alignment sensitivity of SUVIT-OTA as contrasted with Hinode OTA.
Component
M1

Despace

M2

Decenter
Tilt
Despace
Decenter
Tilt

1)

2)
3)
4)

Aberration
DefocusSpherical
Coma
Coma
DefocusSpherical
Coma
Coma

SUVIT-OTA1)(SOLAR-C)
3.4 µm500 µm
48.5 µm2)
4.3 arcsec3)
3.3 µm500 µm
48.5 µm4)
21 arcsec

SOT-OTA1)(HINODE)
8.3 µm
2750 µm
165 µm28.8 arcsec
7.6 µm
2800 µm
165 µm111 arcsec

The positional and angular sensitivities to the mis-alignment of M1 and M2 are evaluated in
the condition that a Strehl ratio is reduced to 0.8 from 1.0 at 633 nm (He-Ne laser) at the field
center.
The coma aberration can be compensated by M2 tilt of 0.189°.
The coma aberration can be compensated by M2 tilt of 0.168°.
The coma aberration can be compensated by M2 tilt of 0.20°.

Mounting of M1 is one of the most critical parts in SUVIT-OTA. The primary mirror
is supported by three stress-free mounting mechanisms seated on the CFRP bottom
panel, interfaced with three superinvar pads bonded on the side of the mirror (Fig. 2.13).
The pad interface of the mounting mechanism provides a kinematic mount for M1. The
pad interface thus avoids stresses to the mirror resulting from dimensional errors inmachining or temperature change. In order to make it sure, the mounting mechanisms
are to be used in the wavefront error measurements of M1 and M2 as noted above. The
only significant surface error of the primary mirror is caused by the difference of CTE
between the superinvar pads and the ULE, which constrains the best-performance
temperature range of the primary mirror to be between -15 to 55℃(TBD).

2-17

M1

M1 support mechanism
Fig. 2.13. Hinode OTA M1 and its supporting mechanism.

M2 is supported by a superinvar tripod-shaped pad of stress relief spring legs that are
glued at the backside of M2. The surface error of the secondary mirror can be again
caused by the CTE difference between the superinvar pad and the ULE, and it
constrains the best-performance temperature range between 0 to 40℃ (TBD). Since thepositional tolerance of M2 is tighter than that of Hinode OTA, it is demanded to havein-flight alignment mechanism of M2 by occasional commanding from the ground. As
shown in Table 2.6, the positional tolerances of M1 and M2 are much relaxed if M2 tilt
adjustments are possible. We expect the mis-alignment of M2 does not occur frequently
but does in post-launch and seasonal intervals. The amount of M2 tilt mis-alignment canbe derived by analyzing a coma aberration of a wavefront error using a phase diversity
method in one of the focal plane instruments.

2.5.4. Thermal design
Solar illumination of about 2 kW is inevitably impinged onto the primary mirror at
the bottom of the SUVIT-OTA during solar observations from its aperture. It is criticallyimportant to make a thermal design to dump such a large heat load to space andmaintain optical and structure components within allowable temperature ranges withsmall temperature fluctuation. In order to identify technical difficulties in the thermal
design of the 1.5mϕ telescope, we investigate thermal environment of the telescopeusing a scaled-up thermal model (Fig. 2.15) of Hinode SOT-OTA except higher solar
absorption coefficient α of M1 and M2 with taking into account the lower reflectivity of
the aluminum coating (α ~ 0.1 for Al coating while α ~ 0.05 for Ag coating at
begin-of-life (BOL), see Table 2.7). The thermal design concept of Hinode SOT-OTA issummarized as follows:
1) Most of the incident energy is reflected back by M1 and dumped out to space by the
2-18

HDM through the heat dump
window on the side of thetelescope.
2) The sunshade has an opticalsolar reflector facing the Sun
to keep it cold, while theupper area of the shield tube
is used as a radiator.
3) The solar heat absorbed in
M1 is radiatively transmitted
to a telescope lower tube fromits side and from a bottom
cooling plate just beneath M1.
4) The solar heat absorbed by
M2 is radiatively transmitted
to the radiator from its
backside.
5) The heat absorbed by the 2FS
and CMU is conductivelytransferred to a mirror celland also emitted out through Fig. 2.14. Thermal design concept in the Hinode OTA.
their housings, and is finallyradiatively transmitted to thelower tube.
6) The heat absorbed by the HDM is conductively transferred to the cylindrical structure
supporting the HDM and outer spiders, and then radiatively transferred to the radiatoron the shield tube and space through the heat dump window.
7) The heat of the lower tube and the shield tube is radiatively emitted directly to thespace through the entrance aperture and indirectly via the radiator of the sunshade
and upper shield tube.
8) The OTA is thermally insulated from the spacecraft. The OTA is physicallyconnected to an optical bench unit (OBU) only by mounting legs of low-thermallyconductive titanium and is radiatively de-coupled from the OBU.
Table 2.7. Solar light absorption coefficient and absorbed energy at begin-of-life (BOL).
Components
M1
HDM
M2
CMU
Total

Absorption coeff.
α
0.12
0.1
0.12
0.11
---

Incident radiation(W)
2074.91830.1
76.2
19.4
----

2-19

Absorbed radiation
(W)
244.8
175.4
9.0
2.1
431.4

300

1940

Φ2233
shield tube
M1-M2
2800

HDM
cylinder

Opticalbench unit

2931
2296

Φ60
truss
pipe

1680
1680

Φ2120
Lower tube

300

Fig. 2.15. Thermal math model for SUVIT-OTA.

Predicted temperatures of optical components at BOL are given in Table 2.8, where
the orbit of the spacecraft is either the sun synchronous polar orbit (SSO) or thegeo-synchronous orbit (GEO). The base model is simply a scaled-up model of Hinode
SOT-OTA, and we find that it gives unacceptably high temperatures both for the
primary mirror and HDM, especially in the case of SSO. When we increase the area ofthe radiator on the shield tube (Mod-1) and increase the conductivity between HDM andthe HDM cylinder and increase the area of the HDM cylinder (Mod-2), it is possible to
make the temperatures marginally within the acceptable range even in the case of SSO.
Table 2.8. Predicted temperatures of the components at BOL.
Components
M1 surface
M2 surface
HDM surface
CMU

Base model
SSO
GEO
61
49
-2 ~ -7-24
109 ~ 111100
34 ~ 37
22 ~ 24

Mod-1
SSO
52
-12 ~ -8
103 ~ 104
24 ~ 27

Mod-2
SSO
52
-12 ~ -7
60 ~ 61
24 ~ 26

It turned out, however, that those modifications still cannot lower the temperature of
M1, especially when the solar absorption coefficient αis degraded at end-of-life (EOL)
and the heat absorption at M1 increases by about 100W (Table 2.9). To lower thetemperature in this case, more drastic modifications are necessary, in which an
additional heat dump system from M1 (e.g. change the bottom of OBU to a radiator) is
accommodated (Mod-3). The prediction based on the modification indicates that we canobtain the acceptable temperatures lower than 50 ℃ even at EOL if the heat load of
about 100W at M1 is dumped at the bottom of the telescope and is transferred toward
the spacecraft bus (Table 2.10). It is described in section 5.5 that there is a feasible
concept for transferring the heat load from the telescope to the spacecraft bus, and fordumping it from a radiator on the side surface of the bus.
It is clear from the thermal model investigation that GEO has an advantage to have

2-20

the moderate temperatures and lower temperature fluctuation because less affected byinfrared radiation from the earth. Another possible solution for lowering thetemperatures is to employ a mirror coating of high reflectivity in UV and low solar heatabsorption like the silver coating (see 2.5.2), which is now under investigation.
Table 2.9. Solar light absorption coefficient and absorbed energy at end-of-life (EOL).
Components
M1
HDM
M2
CMU
Total

Absorption coeff.
α
0.17
0.1
0.17
0.11
---

Incident radiation(W)
2074.91726.3
71.9
18.6
----

Absorbed radiation
(W)
348.6
165.4
12.1
2.0
528.2

Table 2.10. Predicted temperatures of the components at EOL.
Components
M1 surface
M2 surface
HDM surface
CMU

Base model + Mod-2
SSO
GEO
80
70
0~4
-17 ~ -16
63 ~ 6451
41
29

Base model + Mod-2 + Mod-3
SSO
GEO
52
39
-2 ~ 3-18 ~ -17
61 ~ 6247 ~ 48
19
4

2.6. Focal plane instruments
2.6.1. Interface with the telescope
The collimated beam from CMU is spitted by a 50:50 beam splitter (BS) inside theOTA, and is fed into the two focal plane packages (Fig. 2.16). The afocal interface
between the telescope and the focal plane instruments helps to relax the positional
tolerance in the alignment of the packages. The third mirror of CMU can be used as a
tip-tilt mirror (TM) of the image stabilization system, and works to remove the attitude
jitter of the spacecraft in both the beams going to the two packages. The polarization
calibration unit (PCU) is to be inserted between the 2FS and CMU to produce specificpolarization states for polarization calibration, where the polarization modulators are
installed in each focal plane packages as is shown in Fig. 2.6.

2-21

Fig. 2.16. Concept of the beam splitting inside OTA. The beam splitter is located behind the CMU,
and feeds the collimated beam into the two focal plane packages.

2.6.2. Spectro-polarimeter
The spectrograph package consists of the spectro-polarimeter (SP) and the correlation
tracker (CT). The SP is an instrument for obtaining full Stokes parameters of spectralline profiles, and provides precise polarimetric measurements of the chromospheric andphotospheric lines to diagnose magnetic fields in the solar surface. The preliminary
concept of the spectrograph is to employ a Littrow spectrograph configuration similar
with the spectro-polarimeter (SP) of HINODE SOT. The optical configuration of thespectro-polarimeter is shown in Fig. 2.17 (top). The basic parameters are summarized inTable 2.11. It consists of an Echelle grating and an offset-aspheric mirror for collimating
and reimaging the beam before and after the grating. In the current design, the spectrum
bands around He I 1083 nm, Ca II 854 nm, and Mg II h/k 280 nm can be observed
simultaneously in the same diffraction angle with different orders (Fig. 2.17 bottom).
The observing bands include not only chromospheric lines but some photospheric lines
in the observing wavelengths (e.g. Si I at 1082.7 nm) allowing simultaneous observationof both the photosphere and the chromosphere. Is to be switched by changing band-passfilters. Scanning over the FOV across the slit is done with a scan mirror located in frontof the slit. In order to minimize scanning duration, a multi-slit configuration isconsidered, in which two or three slits are located with 60’’ – 80’’ separation to get
spectra at multiple locations simultaneously (Fig. 2.17 bottom). A Wollaston prism isinstalled in front of the detector to measure the orthogonal polarization states
simultaneously. In order to observe the He I line at 1083 nm, it is required to use an
infrared detector instead of CCD. But infrared detectors generally require lowtemperatures environment to suppress noises due to dark currents. We are now
investigating a possible usage of an HgCdTe detector with a 1.7 µm cut-off which canbe operated with relatively higher temperatures (T~200K). The detector is also sensitive
to 854 nm with a good efficiency. In order to observe Mg II h/k lines at 280 nm in SP,
getting co-focus at the slit is optically difficult because the wavelength is very different
from 1.1 µm and 854 nm. It is under investigation to make an optical design allowing
observations at 280 nm in SP.
The precise polarimetric measurements for diagnostics of chromospheric fields
require 10 – 20 sec integration at each slit position to achieve the 10-4 polarimetricsensitivity. In addition to the deep polarimetric observation, the instrument is to be

2-22

designed to support rapid scanning to make spectroscopic observations for study ofchromospheric dynamics and a spectro-polarimetric observations for photosphericmagnetic fields over the FOV with 0.1 sec – 1 sec integration at each slit position. Basicobserving modes with the instrument are summarized in Table 2.12.
Table 2.11. Basic parameters of the spectrograph.
Type
Entrance pupil
Focal length @ slit
Detector size
Pixel scale
FOV along slit
Blaze angle
GrooveSlit width
Sampling dispersion
Wavelength bands
Spectrum lines
Order
Wavelength sampling

Littrow spectrograph
60 mm
1650 mm (F/27.5)
2048 x 2048, 18µm/pixel
0.09 arcsec /pixel
184 arcsec
63.4°
110 / mm
TBD µmλ/∆λ=2.5 x105
1083 nm
He I
15
43 mÅ

854 nm
Ca II
19
34 mÅ

280 nm
Mg II
58
11 mÅ

Fig. 2.17. Optical configuration of the spectro-polarimeter (top), which allows simultaneous
observation of the three wavelength bands at around 1083 nm, 854nm, and 280 nm at the same
diffraction angle (bottom). The figure shows optical rays at each spectral band when three slits arelocated at the image plane.

2-23

Table 2.12. Basic observing modes of SP in SUVIT.
Rapid mode

Normal mode

Deep mode

High-speed spectroscopic observation for studies of dynamics.
- Spectral lines: Mg II h/k 280nm, Ca II 854 nm- Integration time: ~0.1 sec at each slit position.
- Duration: ~2 minutes to scan 100 arcsec FOV.
Spectro-polarimetric observation for study of photospheric magnetic
fields and chromospheric dynamics.
- Spectral lines: He I 1083 nm, Si I 1082.7 nm, Ca II 854 nm- Integration time: ~1 sec at each slit position.
- Duration: ~20 minutes to scan 100 arcsec FOV.
Spectro-polarimetric observation for study of chromospheric magnetic
fields.
- Spectral lines: He I 1083 nm, Si I 1082.7 nm, Ca II 854 nm- Integration time: 10 – 20 sec at each slit position.
- Duration: 20 – 40 minutes to scan 10 arcsec FOV.
(Sparse scanning is used to map wide FOV.)

2.6.3. Broad-band filtergraph
The purpose of the broad-band filtergraph (BF) is to get a monochromatic image ofthe solar photosphere and chromosphere using interference filters at the best possible
spatial and temporal resolution. The optical design is very simple and is shown in Fig.
2.18, in which a filter wheel and a shutter are located in front of the detector. The basic
optical parameters and candidates of the filter pass-bands are shown in Table 2.13. A
key for developing the channel is to make an optical design allowing co-focus at around280 nm and longer than 400 nm.

Fig. 2.18. Optical configuration of the broad-band imager.
Table 2.13. Basic parameters of the broad-band imager.
Entrance pupil
Focal length
Detector size
Pixel scale
FOV

60 mm
2475 mm (F/41.3)
4096 x 4096, 12µm/pixel
0.04 arcsec /pixel
164 arcsec x 164 arcsec

Pass-bands
Spectrum lines

279 nmMg II k/h
2-24

380 nm
CN band

393 nmCa II K

TBD
Contin.

2.6.4. Narrow-band filtergraph
The narrow-band filtergraph (NF) is an important instrument to get monochromatic
imaging in chromospheric and photospheric lines, and to make Dopplergram andmagnetogram observations with covering relatively wide field of view. A tunable filter(TF) is a key component in the narrowband filtergraph. There are two possible types of
TF for SUVIT, (1) Lyot filter and (2) Fabry-Perot filter. The Lyot filter is anarrow-band filter using a birefringence of calcite blocks, where peak wavelengths of its
transmission can be tuned with either rotating wave plates or liquid crystal variable
retarders (LCVRs). There are two possible ways in the Fabry-Perot filter. One is to use
two glass plates and adjust the gap between them to tune the central wavelength of thetransmission (air space). The other is to use a plate made of ferroelectric crystal(LiNbO3) and apply high voltages to tune the wavelength. There are advantages anddisadvantages among them, which are summarized in Table 2.14, and a trade-off study
is now under going.
Table 2.14. Comparison between Lyot filter and Fabry Perot filter for SUVIT.
Lyot filter
Fabry Perot
Air space
LiNbO3
Speed of incident beamF ~ 40F ~ 200
F~90
Necessary diameter
D = 1.5m, FOV=2’ϕ~36 mm~180 mm~81 mmD = 1.5m, FOV=3’ϕ~54 mm~270 mm~120 mmTransmission
~5 %
~70 %
~70 %
Simultaneous 2-pol.
impossible
possiblepossible
Simultaneous multi λ(in principle possible)
impossible
possible (y-cut)
Structurecomplexhigh accuracy
simpleOil
necessary
free
freeControl device
rot. WP or LCVR
Piezo
LiNbO3
Past experience
HINODE SOTSOHO LASCO/C1
SUNRISE IMAX
Advantage- Smaller aperture- Fast tuning- Light and simple- Tolerant- High throughput- Smaller aperture- Large FSR(than air speace FP)- Stable flat field
Disadvantage- Many pol. elements- Big and heavy- High voltage- Need oil- High precision- Slow tunig- Low throughput- Unstable flat field- Small FSR- Long- Small FSR- Unstable flat field- Vulnerability

A possible optical design is shown in Fig. 2.19, where we use a Lyot filter as TF, andlocate it in a telecentric beam of F/40 to make a uniform wavelength over the field ofview. Several wavelength pass-bands can be switched by changing blocking filters
located in front of TF. The reimaging lens is possibly shared with the broad-band
filtergraph in the imager package. The basic parameters of the optical configuration aresummarized in Table 2.15. Requirements on TF are summarized in Table 2.16. The FOV
available with NF depends on how big the TF diameter can be fabricated. When the

2-25

diameter of the Lyot filter is 50 mmϕ at the telecentric focus of F/40, the FOV is to belimited within 120” x 120”. In the case of the air-space Fabry-Perot filter at the
telecentric focus of F/200, the FOV is to be limited within 90” x 90” when the diameterof the TF is 180 mmϕ.

Fig. 2.19. Optical configuration of the narrow-band imager in the case when the Lyot filter is locatedat the telecentric beam of F/40.
Table 2.15. Basic parameters of the narrow-band imager.
Entrance pupil
Telecentric beam
Focal length
Detector size
Pixel scale
FOV

Pass-bands
Spectrum lines

60 mm~F/40 (in the case of a Lyot filter)~F/200 (in the case of a air-gap Fabry-Perot)
1980 mm (F/33)
4096 x 4096, 12µm/pixel
0.05 arcsec /pixel
205 arcsec x 205 arcsec (1)120 arcsec x 120 arcsec (2)90 arcsec x 90 arcsec (3)517 nm
Mg I b589 nmNa I D630 nmFe I

656 nmHIα854 nmCa II(1) The optical configuration allows us to observe the FOV if the TF is big enough to cover the FOV.
(2) The FOV is available when the diameter of the Lyot filter is 50 mm ϕ at the F/40 telecentric focus.
(3) The FOV is available when the diameter of the Fabry-Perot filter is 180 mm ϕ at the F/200 telecentric focus.

Table 2.16. Requirements on TF.
wavelength range
500 – 870 nm
band width~100 mÅ
Strehl ratio>0.9
FOV~200 arcsec
free spectral range>3Å
tuning range±3 Åtuning speed<5 s to scan FSR
tuning resolution<5 mÅ
repeatability<2 mÅ
uniformity
wavelength
5 mÅ (TBD)
transmission
5%
Stability
wavelength
5 mÅ /daytransmission
1% /dayParastic light
<2 %
Ghost
<1 %

2-26

2.7. Image stabilization system
In order to achieve high spatial resolution and precise polarimetric measurements in
SUVIT of SOLAR-C, it is essential to achieve stable pointing of the telescope for solar
images exposed on the detectors (see section 5.4 for more details). The image
stabilization system (ISS) is to be implemented inside the telescope to reduce image
displacements caused by attitude jitters of the spacecraft and disturbances generated bymoving components (e.g. filter wheels, shutters, etc.). ISS consists of a correlation
tracker (CT), which is installed in the spectrograph package and calculates image
displacements by correlation tracking of solar granules, and a tip-tilt mirror (TM) tocompensate the displacements, and servo control electronics. One of the mirrors in thecollimator mirror unit (CMU) is to be used as the TM inside the SUVIT-OTA while theCT is to be installed in one of the focal plane packages, likely in the spectrographpackage, because it is required to control behavior of the image stabilization system incoordination with slit-scanning observations of the spectro-polarimeter.
In HINODE SOT, we have developed ISS to remove residual pointing errors that theattitude control system cannot suppress in the frequency range lower than 14 Hz (Fig.
2.20 left), and achieved the pointing stability of 0.03 arcsec (3σ) as in-flight
performance (Shimizu et al. 2008). It is expected to achieve better stability than 0.015
arcsec (3σ) in SOLAR-C SUVIT for the spatial resolution better than 0.1” and for thehigh polarimetric accuracy because pointing fluctuation during the polarization
modulation introduces artificial polarization signals. High frequency disturbances
induced by moving components and structural resonances are to be partially suppressedby extending the control bandwidth of ISS from 14Hz to about 100Hz (Table 2.17). It isa key to make use of a high-speed (faster than 5000 frames per second) camera with a
CMOS device for correlation tracking of granule patterns. The simulated performanceof ISS is shown in Fig. 2.20 (right).

Fig. 2.20. Power spectrum densities of the pointing error signals measured with the CT in SOT (left).
Simulated reduction performance of pointing error against input disturbances in SOLAR-C with theparameters given in Table 2.17 (right).

2-27

Table 2.17. Basic parameters of the servo-loop performance of ISS.
HINODE SOTSOLAR-C SUVIT
Cross-over frequency
14 Hz> 100 Hz
-4
Time const of TM drive
7.24 x 10 sec< 5 x 10-4 sec-3
CT signal delay time
3.2 x 10 sec< 2 x 10-4 sec

2.8. Contamination control
The contamination of the optical components becomes critical to carry out anobservation in the UV wavelength range shorter than 400 nm. In the HINODE/SOT,
gradual degradation of the throughput has been observed through its mission period.
The degradation is supposed to be because of accumulation of molecular contaminants,
which are out-gassed from organic components inside the telescope, on the surfaces ofthe primary mirror or the collimator lens unit. The decrease of the throughput is about
40% at 388 nm and about 50% at 396 nm 3 years after the launch. The corresponding
throughput degradation at 280 nm is roughly estimated as ~20% if we scale anabsorption coefficient of the contaminants to the shorter wavelength. In order to keep
higher throughput in UV through the mission period, it is required to have more careful
control of contamination environment in designing and developing the telescope. One
possibility is to extend the period of bake-out of components in a vacuum chamber.

2.9. Integration and test plan
2.9.1. Integration and initial M1-M2 alignment
The SUVIT-OTA optical components are assembled in a dedicated integration andtest tower structure having a dummy optical bench unit (OBU) that simulates amounting interface of the spacecraft. The tower structure is capable of tilting the OTA
from a vertical to horizontal configuration as well as rotating around the optical axis
(Fig. 2.21). The integration is performed in the vertical configuration. After the CFRP
truss structure is installed on the dummy OBU, M1 is installed on the mirror cell and its
optical axis is aligned to the center of the mirror cell and the vertical line. Next, the ringplate assembled with the M2 and HDM units is installed on the top ring.
Then the OTA is tilted to the horizontal configuration. De-center, tilt and de-space ofM2 with respect to M1 is adjusted using a Shack-Hartmann sensor attached to themirror cell. The sensor can yield real-time coefficients of low-order Zernike
polynomials such as focus (A20), primary astigmatism (A22, B22), coma (A31, B31)
and spherical (A40) aberration. A high-precision folding flat mirror is set in front of the
OTA entrance aperture. The Shack-Hartmann sensor is focused at the center of thesecondary field stop, and the defocus and coma aberration is eliminated by adjusting
positions and tilt of M2.

2-28

Fig. 2.21. OTA vertical configuration for telescope integration (left) and horizontal configuration for
initial M1-M2 alignment with a Shack-Hartmann sensor and for optical performance test with aninterferometer (right). By rotating the OTA around the optical axis in the horizontal configuration, it
is possible to measure the wavefront error with canceling the gravity effect.

2.9.2. Zero-gravity optical performance test
The wavefront error of the OTA is measured in the horizontal configuration (+1 G
condition) with an interferometer after integrating the CMU to the OTA. Theinterferometer sent a collimated beam into the OTA from the exit pupil. With the folding
flat ahead the entrance pupil, the double-path wavefront error of the OTA is obtained.
Next, rotating the OTA 180 degrees around the optical axis and setting theinterferometer in the opposite side, the measurements for -1 G condition is performed.
Averaging the -1 G and the +1 G wavefront errors can canceled out the effect of gravity.
We can thus measure the wavefront error in the zero-gravity condition.
2.9.3. Opto-thermal test
The thermal model predicts that the OTA on orbit will inevitably have a large
temperature gradient along the optical axis; the bottom part of telescope including M1
will reach a temperature from room temperature to higher than 50 ℃, while the upper
structures supporting M2 will reach a temperature of 0 to 30 (TBD), depending on
thermal conditions and its life phase. It is absolutely necessary, therefore, to verify the
optical performance in such a temperature distribution even though the opto-mechanicaldesign predicts that those temperatures are acceptable.
We will conduct opto-thermal testing of the OTA using a vacuum chamber dedicatedfor space environment tests of optical sensors. The configuration of the test setup is
shown in Fig. 2.22. The OTA is mounted in a vacuum chamber in a horizontal
configuration with a dummy OBU made of superinvar, and the folding flat mirror with atilt and rotational stage is installed at one end of the chamber. Mounted on the side wallof the chamber, the interferometer looks into the OTA through an optical quality
window. The vacuum chamber has upper and lower shrouds surrounding the entire OTA.
By controlling temperatures of the shrouds, we can change the temperatures of upperand lower parts of the OTA independently, while the temperature of the folding flat
mirror is always kept about the room temperature, using another temperature-controlled
shroud.
2-29

interferometer

Fig. 2.22. Configuration to the OTA opto-thermal test.

2-30

Appendix A. Diagnostics of dynamics and magnetic fields in the chromosphereA.1. Limitations in the Zeeman diagnostics
The only way to obtain quantitative information on magnetic fields in the solar
atmosphere is via the measurement and physical interpretation of the emergent spectralline polarization. The approach commonly used is to analyze polarized line profiles dueto the Zeeman effect. The polarization of the Zeeman effect is due to splitting of aspectral line in wavelength. An observational signature of the circular polarization
produced by the longitudinal Zeeman effect is an anti-symmetric profile whose
amplitude scales with the ratio, R, between the Zeeman splitting and the Doppler
broadened line width.
R=

1.4 × 10 −7 geff λB
1.663 × 10 −2 T / µ + ξ 2(1)

where geff is the effective Lande factor of the transition, µ is the atomic weight of theatom under consideration and λ is in Å, B in gauss, T in K and the non-thermal velocity
ξ in km s−1. The linear polarization amplitudes of the transverse Zeeman effect scaleinstead as R2. The Zeeman effect produces significant amplitude of polarization when
the line splitting is comparable with the line width, i.e. R~1, which is the case ofphotospheric lines. On the other hand, the Zeeman splitting is typically much smaller
than the line width (i.e. R<<1) at the major chromospheric lines, because of relatively
weaker field strength and larger line width in the chromosphere. This makes
polarization amplitude smaller there especially in the linear polarization to obtain thetransverse component of magnetic fields (Table 2.1 and Fig. 2.23). This is why we needhigh polarimetric sensitivity to diagnose magnetic field in the chromosphere.
A.2. Prospects of the Hanle diagnostics
There is another physical mechanism by means of which the magnetic fields of thesolar atmosphere makes polarization of the emergent spectral line radiation: the Hanle
effect. Anisotropic radiation pumping processes produce atomic level polarization (i.e.,
population imbalances and quantum coherences among the magnetic sublevels
pertaining to any given degenerate energy level). The Hanle effect can be defined as any
modification of the atomic level polarization due to the presence of a magnetic field.
The Hanle effect makes changes in the polarized line profiles at magnetic strengths
between ~0.1BH and ~10BH (Fig. 2.23), where the critical Hanle field strength (BH, ingauss) is that for which the Zeeman splitting of the energy level under consideration issimilar to its natural width:(2)
B H = 1.137 × 10 −7 t life gJ
with gJ the level’s Landé factor and tlife its radiative lifetime in seconds. The critical
Hanle field strength is typically from 1mG to 10 G dependent on the spectrum lines.
This means that the Hanle effect is sensitive to weaker magnetic fields than the Zeeman
effect (from at least 1 mG to 100 G), regardless of how large the line width due toDoppler broadening is. When the magnetic field is much stronger than the critical
Hanle strength, the polarization amplitude does not depend on the field strength- a1 -

(so-called Hanle saturation regime). Note that in the saturation regime the linear
polarization signals created by atomic level polarization are still sensitive to themagnetic field orientation. When the field strength is strong enough, the linear
polarization of chromospheric lines is generally due to the Hanle and transverse Zeeman
effect (Fig. 2.23). Note also the circular polarization is always dominated by thelongitudinal Zeeman effect, and the amplitude of the circular polarization provides
information of the longitudinal field. The Hanle effect as a diagnostic tool is not limited
to a narrow solar limb zone. In particular, in the forward scattering geometry of a solar
disk center observation, the Hanle effect creates measurable linear polarization in thepresence of an inclined magnetic field (Trujillo Bueno et al. 2002).

Fig. 2.23. (Left) Polarized spectrum line profiles at the He I 10830 Å for several field strengths (ingauss). The critical Hanle field strength BH is about 0.8 gauss at the line. (Right) Polarization
amplitude induced by joint action of the Zeeman and Hanle effect at He I 10830 Å as a function of
the field strength (right). This is the case of forward scattering at the disk center and when themagnetic field is purely horizontal (for the linear polarization) or purely vertical (for the circularpolarization). The Hanle effect allows us to diagnose transverse fields weaker than 100 gauss, which
is not accessible with the transverse Zeeman effect only.

A.3. Inversion of chromospheric lines
The plasma velocities, temperatures, magnetic field vectors, and other interesting
physical conditions in the solar atmosphere are not directly observable. In order toextract quantitative information on such conditions we need to make use of ourknowledge on how the properties of the light are transformed as it propagates. A typicalapproach involves a non-linear fitting of observations of intensity and polarized
spectrum profiles with synthetic ones obtained by solving the forward problem (i.e.
radiative transfer) in a model atmosphere involving some unknown parameters. Because
the chromosphere consists of complex fine structures organized by magnetic fields,- a2 -

interpretation of chromospheric spectral lines requires simplification to restrict the
number of unknown parameters that an observed line profile depends on.
A.3.1. Obtaining velocities and plasma parameters
One of the traditional methods to retrieve plasma parameters of chromospheric finestructures is so-called cloud modeling which describes the transfer of radiation throughstructures located higher up from the solar photosphere resembling clouds on earth’s sky.
Such cloud-like structures, when observed from above, absorb the radiation comingfrom below. The intensity line profiles mostly depends on transparency of the cloud tothe incident radiation and also on the physical parameters in the cloud. Thus, in its mostbasic form cloud modeling delivers the four parameters, cloud optical thickness, cloudsource function, cloud Dopplershift (i.e. velocity), and cloud Dopplerwidth, all four fora homogeneous cloud irradiated from below (see review by Tziotziou et al. 2007). The
cloud source function can be converted into density and temperature through non-LTE
calculation (e.g. Beckers 1972, Tsiropoula and Schmieder 1997) including static
ionization equilibrium. Recent MHD simulations have shown that in chromospherestructures that suffer repetitive episodic heating, the actual populations of the atomic
levels can be far from time-independent statistical equilibrium (Leenaarts et al. 2007).
It is important to refine the cloud modeling to include time-dependent ionization as wellas to compare data directly with numerical spectral line synthesis employing MHD
simulations of time-dependent chromospheric structures.
A.3.2. Obtaining magnetic fields
Inversions of polarized Stokes profiles for photospheric lines have been developed
and used routinely for over 20 years now to produce magnetic field maps. The simplest
one is to use the Milne-Eddington method (e.g. Skumanich and Lites 1987) employing
simplified assumption that all the physical parameters are constant in the line-formation
region (except for the source function which varies linearly with optical depth) and has
the advantage that the forward problem has an analytical solution, which makes theinversion robust and fast. The second method, the Stokes inversion based on response
function (SIR, Ruiz Cobo and del Toro Iniesta 1992), is based on a full solution of thepolarized LTE transfer in an arbitrarily one-dimensional atmosphere. It has theadvantage that it is possible to determine the variation of physical parameters along theline of sight.
Similar tools for chromospheric inversions already exist and significant experiencehas been accumulated for nearly a decade now. There are several reasons why the
analysis of chromospheric lines is more challenging. The biggest issue is that theradiative transfer in the chromosphere is non-LTE and there is not a universal solution.
One needs to employ different techniques depending on the particular physical scenario
under study (whether the polarization is originated by Zeeman effect or scattering, etc).
They fall roughly into two different approaches. The He I triplet at 1083 nm forms overa very thin atmospheric slab because it forms by recombination of He II ions created by
UV coronal illumination. The line formation allows us to use the Milne-Eddington typeof approach as a valid approximation to derive accurate magnetic fields in the Zeeman
regime (Lagg et al. 2004). The inversion code HAZEL (Hanle and Zeeman Light,- a3 -

Asensio Ramos, Trujillo Bueno & Landi Degl’Innocenti 2008) is able to tackle the fullproblem including the Zeeman and Hanle effects self-consistently in such a slab (Fig.
2.24).

Fig. 2.24. Example of a spectro-polarimetric observation of He I 10830 Å line and non-linear fittingof the observed line profiles with HAZEL.

Another alternative is the generalization of SIR to non-LTE atmosphere(Socas-Navarro et al. 1998). In the current code, the polarization generation is restricted
to the Zeeman regime only but the line formation is treated fully in non-LTE, with theassumptions of instantaneous ionization equilibrium and complete angle and frequency
redistribution (CRD). The CRD approximation is valid for some lines such as those in
the Ca II infrared triplet at around 850 nm (Uitenbroek 1989). The non-LTE inversion
was successfully applied to spectro-polarimetric observations in the Ca II infrared triplet
to get three-dimensional magnetic structures around a sunspot (Socas-Navarro et al.
2005). The validity of the inversion and the possible impact of ionization history andthree-dimensional radiative transfer are currently been studied using numerical MHD
simulations of the chromosphere (de la Cruz Rodriguez et al. 2010, Fig. 2.25) thoughthese approximations affect only the determination of thermodynamical parameters but
not the magnetic field vector and the line-of-sight velocity.

Fig. 2.25. Test of the non-LTE inversion in the Ca II 854 nm line using the numerical MHD
simulation of the chromosphere (see A.4), where top panels show parameters from the simulation
and bottom panels show parameters obtained with the inversion of the synthesized line profiles.

Because the non-LTE inversion requires lots of computational power, it is underinvestigation to apply pattern recognition techniques such as principal component- a4 -

analysis (PCA, Casini et al. 2005) and artificial neural network (ANN, Socas-Navarro2005) to the inversion of the chromospheric lines. The general idea behind thepattern-recognition approach is to create a sufficiently dense and uniform database ofmagnetic and atmospheric models that are relevant for the line formation characteristics
in an observation. The computationally intensive forward-model step is thus pushed
upfront to a one-time calculation of an adequate database of Stokes profiles, while theinversion problem is reduced to the search in that database of the best match to theobserved line profiles.
A.4. Numerical modeling of the chromosphere
In order to fully exploit detailed chromospheric line observations, including theirpolarization measurements, a promising approach would incorporate comparisons ofobserved spectral line properties and dynamic phenomena with those seen in numerical
MHD simulations of the same observables. This requires 3D models of the photosphereand the chromosphere including a realistic description of the important physics. Fromsuch models it is then necessary to calculate the observable spectral lines employing
radiative transfer.
Although there are lots of complexities in the modeling, several codes exist within
the description of single-fluid magnetohydrodynamics. They can treat the detailed
radiative transfer, and their photospheric simulations reproduce detailed observations toremarkable detail (e.g. Stein and Nordlund 2006). In the chromosphere, scattering starts
to play an important role, radiative losses in strong lines from calcium, magnesium andhydrogen, where non-LTE is crucial, and the hydrogen ionization balance is out ofstatistical equilibrium. Recently numerical codes have been developed that can treat allthese effects and simulations have been produced where comparisons with observations
start to become meaningful (Fig. 2.26, e.g. Leenaarts et al. 2009). Efforts are also
underway to go beyond the single-fluid magnetohydrodynamic description of thechromosphere.
Once a realistic simulation exists it is necessary to calculate the observable radiationin order to compare with the observations. It is important that the employed radiativetransfer codes are capable of handling partial frequency redistribution (PRD), which
accounts for the coherence of the incoming and outgoing photons in scattering events instrong chromospheric spectral lines like the Mg II h/k and the Ca II H/K lines. In
addition, to compare with observations of polarization in chromospheric lines, theradiative transfer codes should be able to calculate polarized line profiles from a given
three-dimensional magnetic configuration in an MHD simulation. Several non-LTEtransfer codes already exist that meet many of the requirements, and in the near future
these efforts will be combined.

- a5 -

Fig. 2.26. Snapshot of a 3D radiative MHD simulation including the photosphere, chromosphere,
and corona (courtesy of Carlsson and Hansteen).

- a6 -

Chapter 3: EUV/FUV High-Throughput Spectroscopic Telescope
3.1. Performance requirements for the spectrometer
The Solar-C mission provides an ideal platform for a EUV/FUV high-throughput highresolution spectrometer (may be called EUVS hereafter). Performance requirements are
summarized below for a breakthrough spectroscopic instrument, which can be realized withrecent technology advancements.
1. To operate in both spectroscopic and imaging modes with high temporal resolution(cadence). This is required for quantitatively understanding the physics of dynamical
phenomena. One order of improvement in effective area is required. Thus, thetemporal resolution at a slit position can be shortened down to 0.5 sec for active region
plasma and 1 sec for quiet Sun plasma.
2. To perform EUV/FUV spectroscopy with high spatial resolution. The spectrometer isrequired to resolve structures equal to 0.5” resolution or better. The target is 0.3 arcsec,
which is roughly equivalent to the width of chromospheric spicules observed with
Hinode/SOT.
3. To observe a field of view as large as possible. The best resolution may be restricted
to ~5x5 arcmin2 (a medium size active region), but lower resolution should cover alarger field of view.
4. To perform EUV/FUV spectroscopy with high spectral resolution. The spectrometer isrequired to determine Doppler velocities to an accuracy of ~ 2 km/s from spectral line
shifts and non-thermal motions as small as 20 km/s from line widths.
5. To make measurements with spectral lines from a broad temperature range and withadequate plasma diagnostics covering chromospheric, transition region, low coronaand flare temperatures. The observations must cover all relevant temperatures andstructures in order to track the flow of energy from the chromosphere into the corona.
6. To perform good observations with low-scattering optics. This performance is needed
to explore “faint” reconnection outflow regions, coronal holes, and off-limb
observations.
7. To ensure precise co-alignment with the data from other onboard telescopes. Coalignment must meet the high spatial resolution of the spectrometer.
The performance requirements summarized above are determined from the science goals in
section 1. Table 3.1 summaries instrument performance required for the Solar-C spectrometer,
compared with the previous spectrometers. It is obvious that the Solar-C spectrometer will
have significant improvements on the temporal resolution, spatial resolution, and temperature
coverage. Some of the performance required for the Solar-C spectrometer are discussed below
to show possible discovery spaces and new sciences.

3-1

Table 3.1. Performance of Solar-C spectrometer in comparison with other spectrometers
Temporal resolution fora slit position *1
Spatial resolution

Solar-C spectrometer
0.5 - 10 sec for AR
1 -10 sec for QS~0.3 arcsec (target)~0.16 arcsec (pixel size)
4.1 – 7.2

Hinode/EIS
10 - 20 sec for AR
30 – 60 sec for QS
3-4 arcsec
1arcsec (pixel size)
4.9, 5.6-7.2

SoHO/SUMER
10 – 20 sec for AR line*2
2 arcsec
1arcsec (pixel size)
4.1-5.8

Primary temperature
coverage (log10 K)*1 The values given here are typical exposure duration for multiple numbers of spectral lines covering a wide
temperature range.
*2 SUMER can record spectral lines in only one of spectral windows at the same time. If spectral lines inanother spectral window are measured, the grating is needed to move before the exposure.

3.1.1. Spatial resolution requirement
One of key tasks of EUVS is to precisely measure plasma conditions, such as temperature,density, and velocity of elementary structures of the magnetized atmosphere in all
temperature regions from the chromosphere through transition region and into the corona.
Such measurements made simultaneously with different spectral lines allow us to trace theflow of energy from the chromosphere up into the corona. In the photosphere the energy isobserved to be in plasma structures with a scale length of about 0.3″ or about 200 km. Movies
made in the Ca II H-line from Hinode SOT clearly reveal that the chromosphere is dominatedby a multitude of thin (~200 km), dynamical, jet-like extrusions called spicules. Since SOT
can carry out only imaging observations for the chromosphere without spectroscopic
capabilities, the proposed spectrometer will be the first instrument to diagnose the plasma
properties of these thin, dynamical structures and to follow these structures upwards throughthe atmosphere into the corona. As the magnetic field may expand with height, the overall
size of structures may increase as they rise above the chromosphere/transition region, but thewidth of spicules may be preserved into the transition region and corona. Coronal loops are
well-known coronal magnetic structures, but elementary structures are not yet fully resolvedin the corona with the currently existing imaging telescopes (0.5
″ at the highest resolution).
Spectroscopic measurements in the EUV have suggested that only 10% of the volume ofcoronal loops is filled with hot plasma, when they are observed with the spatial resolution of
2-3 arcsec (pixel size 1 arcsec) (Warren et al. 2008). Thus, measuring the plasma parameterswith 0.3 arcsec, i.e., 10% of the 1 arcsec volume, can fully resolve elementary structures
hidden in the magnetic loops. Understanding the magnetic sub-structure implied by the small
filling factor is crucial for determining energy flow in coronal loops and ultimately forunderstanding the physics of coronal heating.
In addition to its own scientific contribution, the spectrometer observations shall haveclose scientific synergy with the other proposed telescopes. One of the important science
goals is to understand the roles of waves propagating through the atmosphere in heating andactivating dynamical motions in the plasma. It is important to determine whether or not
sufficient energy remains in the waves that penetrate all the way to the corona to make a
substantial contribution to the heating budget of the solar corona and/or to the acceleration of
the solar wind. It is also critical to identify propagating MHD waves and their modes. For thispurpose, observations of magnetic field fluctuations are required in addition to observationsof intensity and velocity fluctuations. The proposed spectrometer can provide the observationof intensity and velocity fluctuations. Magnetic field measurements are available only for the
low atmosphere, i.e., in the photosphere and chromosphere, and they will be provided by the
3-2

UV-Visible-IR telescope with a spatial resolution of″ 0.2
or better. The UV-Visible-IR
telescope is designed to determine the modes and amplitudes of waves in the low atmospherejust after the excitation of the waves, and the proposed spectrometer can continue to trace
their propagation into the transition region and corona. With these coordinated observations,
we will explore for the first time the mode conversion in waves, which is expected in theatmospheric layer where the plasma beta is unity. Since the plasma beta changes from high(β>1) in photosphere and low chromosphere to low (β << 1) in the corona, we can observe
how the waves are evolved as a function of height in the stratified atmosphere. With having
high spatial resolution never achieved so far, we can newly explore behaviors of waves andturbulence in small scale, which are useful information for physically understanding roles ofwaves in the atmosphere. We may have new information concerning turbulent cascade andphase mixing of Alfvenic waves.
Considering these new discovery spaces, the proposed spectrometer should have thecapability of resolving structures with scales smaller than 350 km (0.5″). Ideally, it would bebetter to resolve even smaller structures, say, 0.3″ or ~200 km.
3.1.2. Temporal resolution requirement
Since the structures to be resolved with EUVS are small and dynamic, a high time cadenceis necessary and consequently a high-throughput instrument is required. As a goal of the
instrument design, the effective area of the spectrometer is set to be at least 10 times larger
than previous instruments (EIS onboard Hinode, SUMER and CDS onboard SoHO). Thismeans an improvement of at least an order of magnitude in temporal resolution if observing
with the same sampling scale (~1″ arcsec). There is also an improvement even whenobserving with higher spatial resolution (0.5″ arcsec) than achieved with previous instruments.
In the case of EIS observations, typical integration times of an exposure are 10 to 20 s for
active region emission lines and 30-60 s for the quiet Sun (Note that shorter exposures are
possible in very limited numbers of brighter lines). SUMER observations are carried out with10 to 20 s integration times (shorter exposures would be possible in the brighter lines but the
exposure time is limited by the available telemetry rate). With having high throughput
performance and bright spectral lines, the integration times of the proposed spectrometer canbe shortened down to 1-0.5 s. The sit-and-stare measurements at a slit position provide timeseries of spectral data with temporal resolution of 0.5-1 s. The field of view with 15 arcsec
width can be repeatedly measured with 25-50 sec cadence, if the slit step is 0.3 arcsec. The
spectral maps for 100 arcsec field of view are available every ~200 s, if the slit is stepped per
0.6 arcsec.
As discussed in section 1, numerical simulations suggest that low-frequency (<5mHz)
magnetoacoustic waves may provide a significant source of the energy necessary for
maintaining the chromosphere. Moreover, the outgoing low-frequency (<1mHz) Alfven
waves may provide a major contribution to the heating and acceleration of the plasma as aresult of the nonlinear generation of the compressive waves and shocks in the corona.
Observationally, various kinds of oscillations and waves have been detected with Hinode data.
For example, with analyzing EIS spectral data, upwardly propagating and standing slow
magnetoacoustic waves are reported in active regions (Kitagawa et al. 2010) and coronalholes (Banerjee et al. 2009). When the waves propagates along a coronal loop with the soundspeed (~150 km/s in the corona), the proposed spectrometer can scan the entire coronal loop
about 5 times with 0.6 arcsec step during the time when the waves propagates from one end tothe other end of the loop. Thus, such data will allow us to trace the waves from the excitation
to the dissipation. The speed of Alfven waves is much higher than the sound speed in the

3-3transition region and corona, but the scan speed of the proposed spectrometer can allow us to
explore the properties of propagating Alfven waves, if we choose slit scan step size properly.
Moreover, owing to higher temporal resolution, we will search for higher frequency waves(~1Hz) that cannot be explored with the currently existing spectrometers. If higher frequency
waves are found in the transition region and corona with sufficient flux, high frequency waves
may be a possible candidate for making the bright corona in active regions.
High temporal cadence performance with high spatial resolution has the potential to resolve
fine dynamical structures in the coronal loops. Hinode EIS has revealed excess non-thermalline broadening at the footpoints of coronal loops, which is interpreted as the existence ofhighly blue-shifted small-scale flows at the base of coronal loops (Hara et al. 2008). Nonthermal line broadening is known in spectral lines from the transition region and corona, and
we guess that the broadening is created by turbulent nature in coronal loops. Having short
integration time is definitely useful to distinguish each of turbulent flows, providing new
information on interpreting non-thermal line broadening.
Higher temporal cadence would provide strong tools in exploring coronal holes. Existing
spectrometers need much longer integration times to meaningfully record spectral data incoronal holes, where the FUV and EUV coronal emissions are much fainter than in active
regions. It is important to identify the source region of the fast solar wind in polar coronalholes. Because of very long integration times (several minutes for a slit position and several
hours for mapping a region), it is difficult with existing instruments to confidently identify thesource regions. A substantial improvement in the throughput performance can much reduce
the time needed to spectrally map a region, getting closer to the timescales for the evolutionof the structures. Also, waves with periods from 7 to 20 min have been detected in plumes
and inter-plumes above the limb in the polar regions (DeForest & Gurman 1998; Ofman et al.
2000; Banerjee et al. 2000, 2009). A high throughout performance could allow a search forhigher frequency waves in coronal holes and a search for wave signatures in line widths andvelocity.

3.1.3. Broad temperature coverage
As one of the top-level science goals, we will attempt to understand how the energy istransported through elementary magnetic structures from the chromosphere to the transition
region, and from the transition region into the corona. Furthermore, magnetic energy
dissipation events are well observed over the entire solar atmosphere, not only in the outercorona which is filled with fully ionized, collisionless plasma, but also in the chromospherewhich consists of weakly ionized plasma. In order to capture the entire picture observationally,
it is important to measure the plasma conditions and behaviors in different temperature layers
simultaneously.
SoHO/SUMER records spectral lines at longer wavelengths in the FUV, where many
spectral lines originating from the transition region and chromosphere are available, while
only few coronal lines are observed. In contrast, Hinode EIS observes spectral lines at shorter
wavelength in the EUV, where many spectral lines from the corona and flare plasma are
available but the number of spectral lines is limited for exploring the transition region andchromosphere. Extensive coordinated observations by the two instruments have been obtainedsince the launch of Hinode, providing significant new results. However, we realize that it is
impossible to achieve high accuracy in the co-alignment of the spectral data from SUMER
and EIS. Also, since the spectral data are acquired for a single slit position and it takes a longtime to map a large field-of-view, it is nearly impossible for two different slit spectrographs
to obtain data at the same location on the Sun and at the same time. However, observations at

3-4

the same locations and times are critical for understanding the dynamical behavior andevolution of the plasma in the observed targets, because the timescales for the evolution of
structures are quite short.
Therefore, it is a requirement that the next generation advanced EUV/FUV spectrometer
capture spectral data over the entire temperature domain of the solar atmosphere at the same
locations and times. This performance can much improve and ensure accurate co-alignment
among different spectral lines. With this improvement, we can discuss with high confidence
the relationships among the observational signatures at different temperatures. Elementarymagnetic structures of the order of 0.3″ seen in the chromosphere can then be traced into thetransition region and corona. It is also possible to trace the flow of energy through static
elementary magnetic structures over the solar atmosphere and quantitatively investigate therole of waves as a transporter of energy for heating the corona and energizing the solar wind.
3.1.4. Necessity of low scattering optics
Low-scattering optics should be considered when the spectrometer is designed andfabricated to allow us to accurately record low intensity data. A single-mirror, off-axis
telescope is mandatory for off-limb observations, to avoid stray light from the bright solar
disk. Many important science discoveries may be hidden in low emission measure areas suchas reconnection outflow regions, which are much fainter than flare loops. As another example,
low intensity emission in coronal holes observed off-limb may hold the key for understanding
solar wind acceleration and shock formation in the high corona.
The reconnection region science discussed in section 1.3.4.2 is an excellent example forstressing the importance of low-scattering performance. Probing the reconnection region itself
is not easy because intense flare loops are formed near the reconnection region due tochromospheric evaporation. According to the model calculation given in section 1.3.4.2, thetotal intensity of the reconnection outflow region is about one thousandths of the total
intensity of the bright post-flare loops, when observed in spectral lines from Fe+17 to Fe+23.
This estimate tells us that low-scattering optics are necessary for investigating the flarereconnection outflow region.

3.1.5. Substantial jump expected from the NASA IRIS mission
IRIS (Interface Region Imaging Spectrograph) is a NASA SMEX mission selected forlaunch in December 2012. This spectrograph observes several spectral lines in the 120 to 280
nm range with high spatial resolution (0.3″) and high cadence (1 to 5 sec), to discover how
the outer atmosphere is energized. The spectral lines observed with IRIS are suitable for
exploring the interface between the upper chromosphere and lower transition region, and the
IRIS mission is expected to open a window of discovery into a crucial gap in current solar
observational capabilities.
However, the spectrometer proposed here would provide a substantial step forward from
IRIS observations, because the Solar-C Plan B mission can provide unique observations thathave never been realized, allowing us to progress beyond the chromosphere-transition region
interface into the entire transition region and corona. The proposed spectrometer will be able
to observe spectral lines in the EUV and FUV that seamlessly cover all the temperatures fromthe chromospheric to the corona with high spatial resolution (0.5″ or better) and high cadence(1 to 5 s). Observations of this completeness are necessary for understanding coronal heating
and the formation of the solar wind. In addition, the proposed spectrometer observes several
spectral lines that originate from high temperature plasma in the 106-107 K temperature range,

3-5and these spectral lines are useful for diagnosing flares and active region dynamics in thecorona. Thirdly, and more important, is that EUV and FUV spectroscopic observations canbe carried out simultaneously with quantitative measurements of magnetic fields in thechromosphere and photosphere with 0.1″ to 0.2″, taken by the UV-Visible-IR telescope. This
would be an enormous help in associating an accurate magnetic geometry with the EUV-FUV
spectra. Finally, unlike IRIS, the single-mirror off-axis design of the telescope allows us toobserve faint structures in the corona far from the solar disk.

3.2. Technical feasibility of each key element
The strawman instrumentation draws heavily on current technology under developmentby the international EUV/FUV spectroscopic community. Technical feasibility of each key
technology is briefly described below.

3.2.1. Two element optical layout
In order to improve throughput performance, minimizing the number of optical elements
is essential; the spectrograph consists of two optical elements, i.e., an off-axis parabola
primary mirror and a concave grating. The two-element spectrograph design was firstdeveloped and flown for EIS onboard Hinode and it has since been refined to meet various
requirements for RAISE, VERIS and EUNIS rocket payloads.

3.2.2. The telescope mirror and mirror coating
For off-limb solar observations the single-mirror telescope greatly reduces stray light, as
compared to an on-axis, double-mirror design. In addition, a secondary mirror would not only
reduce throughput by adding another reflection but would also be vulnerable to degradationdue exposure to very high solar flux. Thus the decision is made for a single-mirror telescope.
Increasing the size of the primary mirror gives a significant improvement in throughput. If
the diameter of the mirror can be increased twice, the effective area is improved by a factor 4.
30 cm or larger diameter gives 4 times or larger throughput, as compared with EIS (diameter
15cm).
Driven by requirements for EUV photolithography, EUV optics with very low scatter anddiffraction limited performance are available. For example, the optics developed for the MET
EUV photolithographic system have a slope error of 0.3 to 1.0 microradians, andmicroroughness of 0.1 to 0.2 nm over spatial frequencies from 9 nm to 4 mm (Goldberg et al.
2005).
A single coating on the primary mirror – without division in sections for EUV and FUV
ranges - results in good performance across a broad wavelength range. The suggested coating
is described in Figure 3.1. It consists of a broadband Si/Mo multilayer with a 5 nm B4C
topcoat. The B4C topcoat is semi-transparent in the EUV but gives reasonable reflectance atlonger wavelengths as shown in Figure 3.1 (left). The single coating minimizes risk andcomplexity for the primary mirror associated with dividing the solar illuminated mirror into
segments with different coatings. Recent advances in stabilizing multilayer for higher
temperature performance will be incorporated into the multilayer coating design.

3-6

3.2.3. Grating and grating coating
To realize seamless temperature coverage in the 104 to 107 K range, the spectrometer
needs to cover between 2 and 4 wavelength bandpasses over a range from 10 nm to 120 nm.
To accommodate this broad wavelength range, the grating is divided into two segments. The
ruling density, groove geometry and figure of each segment is optimized for the specificwavelength range. Each grating segment is coated individually. A process is envisioned
where each coated grating segment would be contacted/cemented to a common plate. This
plate would then be installed into the instrument.

Figure 3.1. Reflectivity of the proposed primary mirror coating. The red curve in the left
panel shows the performance of the coating from 50 to 120 nm. The right panel gives the
EUV coating performance.

3.2.4. Solar blind detectors for EUVS
Enhanced technology detectors with larger format and faster readout are underdevelopment for space flight. The intensified (CMOS-APS) active pixel sensor (flown on
NASA’s EUNIS and RAISE rocket flights) has the promise of exceptionally fast image
readout with quick access to specific pixel areas. The intensifier with opaque anode ensuresthat the detector is “solar blind” and focal plane filters, to suppress visible solar radiation, are
not needed. The opaque photocathode on the microchannel plates will ensure very highsensitivity in selected spectral ranges.
At shorter wavelengths (in the EUV), back side thinned, large format, fast readout silicon
based CCD and APS detectors are highly efficient, low-noise, fast-readout detectors. The
backside illuminated detectors will be fitted with focal plane filters to provide rejection of thevisible solar radiation scattered from the grating. The performance of currently available
intensified CCD and EUV backside illuminated CCDs are suitable for the EUVS mission
while similar CMOS-APS sensors are presently under development.

3.2.5. Removal of metal filters and thermal control concept
To achieve low scatter, higher throughput and broad wavelength coverage, the EUVS
strawman design has no prefilters in the front aperture. These prefilters have historically beenused to protect the multilayers from the incident solar radiation. With one solar constantilluminating the primary mirror, the temperature of the mirror is expected to increase to 50 °C
3-7

or 60 °C provided that a passively cooled plate is provided. Similar systems for NEXUS
instrument (proposed as a SMEX) and SUMER (flown on SOHO) have had very reasonableproperties, but a more detailed solution has yet to be found. A heat rejection mirror will beplaced in front of the slit plane to reject the majority of the solar radiation before the slit.

3.3. Strawman Instrument

3.3.1. Optical Design
The EUVS instrument schematic is shown in Figure 3.2. As shown in the schematic,
solar radiation is incident on the EUVS off-axis paraboloid primary mirror (φ~40cm). The
primary mirror focuses an image of the solar disk on the entrance slit and reflective slit jaws.
Solar radiation falling outside the instrument field of view is redirected out the front of the
instrument with a heat rejection mirror following a similar design used for the
SUMER/SOHO and the NEXUS instruments. A slit selection mechanism interchanges andprecisely places various slits in the telescope focal plane.

5m length

Primary Mirror

solar
radiation
Focal Plane Detectors

Entrance Slit

Grating

Figure 3.2. EUVS strawman instrument schematic
Solar radiation passing through the slit is reflected and dispersed by the elliptical,
variable line space grating onto the various focal plane detectors. The grating may be divided
into as many as four separate gratings to cover the various wavelength bands. The figure,ruling, and coating of each grating segment are optimized to meet the requirements of theparticular wavelength range. The spectra are recorded by large format, electronic detectors.

3.3.2. Effective Area
Results from a conservative first order effective area calculation for the strawman
instrument are given in Figure 3.3. The effective area was calculated using the expected
performance of a next-generation, broadband multilayer reflection coating for both mirror andgrating. A groove efficiency of 0.40 and detector efficiency of 0.5 were assumed. Each

3-8channel used 0.25 of the mirror total area. The calculated effective area is a factor of >10
higher than previous instrumentation as shown below. Further optimization and refinement of
this calculation would likely show considerable improvement. In particular, the use of only
two grating sections for the EUV and the FUV ranges will enhance by a factor two thethroughput.

Figure 3.3. Effective area of the EUVS strawman instrument. Previous comparable effective
areas for EIS, CDS and SUMER are plotted for comparison.

3.3.3. Selection of wavelength bands
Four spectral bands were chosen as representative for the purpose of developing thestrawman design. The bandpasses will be further optimized to meet the specific Solar-C
scientific requirements.

3.3.4. Expected count rates and Temperature Coverage
The expected count rates for the emission lines in the four channels have been calculated
and are given in Appendix 3-A. The temperature coverage of the bandpasses is superb as
shown in Figure 3.4. The total counts in each line for quiet Sun, active regions and flares are
given for a spatial sampling of 0.33". Count rates larger than 100 counts/s will allow rasteringof quiet-Sun plasmas with 1-s exposure times in the hydrogen H I and the carbon C III lines.
Spatial binning and slightly longer integrations allow rapid scanning in a wide variety offainter lines (C II, He I, O V, Ne VIII, Fe IX, Fe X, etc.) and will yield extraordinary
3-9spectroscopic images of the quiet Sun. Active region count rates of several hundred counts/s
with 0.33" spatial resolution enable spectroscopic imaging of active region flows anddynamics. The EUVS line list includes high temperature, well isolated emission lines in flares
and active region plasmas.
Figure 3.5 compares the count rates expected for the Solar-C spectrometer with thecurrently operating SoHO/SUMER and Hinode/EIS and planned IRIS. It is obvious that theSolar-C spectrometer collects much more numbers of photons in 1 sec integration time, eventhough the spatial resolution is 0.33 arcsec.

Figure 3.4. Expected Count Rates. The expected count rates are plotted as a function ofionization temperature for a spatial sampling of 0.33". The plots show total counts in each line
plotted at their ionization temperature. Each asterisk corresponds to a specific emission line.
The top panel corresponds to typical quiet Sun; the middle panel shows typical active region
count rates, and the bottom panel shows flare plasma count rates.

3-10

Active Region Plasma
10000

Solar-C/ EUVS
1000

Plan ned IRIS

100

Hinode/ EIS

10
SoHO/ SUMER

1
10 4

10 5
10 6
Temperature (K)10 7

Figure 3.5. Count rates for typical active regions. The Solar-C spectrometer is compared withcurrently operating SUMER and EIS and planned IRIS. The expected count rates are plotted
as a function of ionization temperature for the spatial sampling of each spectometer. Eachasterisk corresponds to a specific emission line observed with the Solar-C spectrometer.
Triangles, diamonds, and squares are for IRIS, SUMER, and EIS, respectively.

3.3.5. Improvements over previous instrumentation
Critical improvements over previous instrumentation are obtained with this strawman
design as follows:•••>18 improvement in spatially resolved area on the solar disk
>10 improvement in effective area
complete temperature coverage (104 – 107 K in active regions with Δ(logT) <0.3)

A summary table of the EUVS scientific characteristics is given in Table 3.2.

3-11

Table 3.2. EUVS scientific summary characteristics.
Plate scale
Simultaneous field of view
Maximum raster width
Wavelength range
Exposure times
Velocity resolution
Temperature coverage~0.16"/pixel (sampling comparable to Hinode SOT)
Slit: 0.16", 0.32", 0.96"; Slot: 10", 40", 200"
Slit/slot length: 200" nominal, >300" extended+/- 75" nominal, +/- 200" extended
four channels (Å):130-250; 510-630; 690-810; 940-1060;
1-5 s with 0.33" spatial sampling
0.1-0.5 s in active regions with 1.0" spatial sampling
Doppler (centroid) shift measurement accuracy <3km/s
Turbulent velocity (width) <10km/s
104 – 107 in active regions with Δ(logT) <0.3

3.4. Strawman design details
Further details of the strawman design are given in Table 3.3. Provisions have been made in
the strawman design for a slit jaw camera imaging system and a possible closed loop servocontrol system using a correlation tracker system and an articulated primary mirror for image
stabilization. Solar image rastering is accomplished by tilting the primary mirror.

Table 3.3. EUVS strawman design – nominal performance
Telescope
Baseline or nominal valueType
off-axis paraboloid mirror, single coating optimized for both long and short bands.
optical properties 40cm usable aperture diameter, 400cm focal length, 30cm off-axis distance, <λ/80rms,<0.5nm microroughness, 19.4 microns/arcsecond plate scale.
Mirror blank
Lightweighted premium grade zerodur blank (or equivalent).
Stabilization<0.03 arcsecond jitter suppression, >5Hz bandwidth possible
Geometric blur ~0.26" at 75" roughly circular field of view,~0.52" at 150" roughly circular field of view
Slit/slot
Baseline or nominal value
assembly
General
6 slits/slot positions, front face with reflective coating, >327" length
selected widths Slits: 0.16", 0.32", 0.96"; Slots: 10", 40", 200"
Slit jaw imaging Baseline or nominal value
GeneralProjecting a high-quality image of the slit jaws onto a CCD camera.
Feeds a CCD camera used for correlation tracker system.
Design char.
field of view: 400" x 400"; plate scale: 0.2"/pixel; CCD format: 2K x 2K.
Filter bandpass Possible selections include: Ly-α line (121.6 nm), Tmin(160 nm), Mg II line (280 nm),
Ca II (393 nm), G-band (430.5 nm), Hα line (656.3 nm), TiO band (705.7 nm).
Correlation
If implemented is expected to be similar to that used on Hinode/SOT bandwidth (10-30
tracker system
Hz). Dedicated small-format CCD camera.
SpectrographBaseline or nominal value
characteristics
GeneralOptimized single reflection EUV spectrograph, elliptical substrate with variable line
spacing as required, grating figure and ruling individually optimized for the
bandpass/quadrant.
Typical blur: <1 pixel over majority of wavelength range for +/- 50" at slit center.
Coating selected and optimized for each quadrant.
Detector optimized for the bandpass/quadrant
EUV1 bandpass detector: 12kx2k mosaic CCD/APS, backside illuminated, 13.5 µm pixel size.
(130-250Å)stray light rejection: thin film metallic filter with supporting mesh.
spatial plate scale: 0.16"/pixel.
spectral plate scale: 0.010 Å/pixel, 17.7 km/s-pixel@170Å.
Grating: m~4.35, elliptical figure, 4000line/mm VLS ruling, -1 order.

3-12

EUV2 bandpass detector: 5kx2k CCD/APS, 17 µm pixel size with MCP intensifier with KBr(510-630Å)
photocathode.
stray light rejection: solar-blind intensifier
spatial plate scale: 0.20"/pixel.
spectral plate scale: 0.026 Å/pixel, 13.7 km/s-pixel@570 Å.
grating: m~4.35, elliptical figure, 2000line/mm VLS ruling, +1st order.
EUV3 bandpass detector: 5kx2k CCD/APS, 17µm pixel size with MCP intensifier with KBr(690-810Å)
photocathode.
stray light rejection: solar-blind intensifier.
spatial plate scale: ~0.20"/pixel
spectral plate scale: ~0.025 Å/pixel, 10.0 km/second-pixel @ 750Å.
grating: m~4.35, elliptical figure, 2000line/mm VLS ruling, -1st order.
EUV4 bandpass detector: 5kx2k CCD/APS, 17 µm pixel size with MCP intensifier with KBr(940-1060Å)
photocathode.
stray light rejection: solar-blind intensifier.
spatial plate scale: ~0.20"/pixel.
spectral plate scale: ~0.024 Å/pixel,7.2 km/s-pixel@1000 Å.
Grating: m~4.35, elliptical figure, 2000line/mm VLS ruling, +1st order.

3.5. Further optimization for better instrument
Considerable optimization activities remain for the spectrograph instrument. Science
goals and associated measurement requirements need further definition. The trades relating to
instrument wavelength ranges need further definition. The number of channels drives
instrument complexity. The wavelength ranges should be placed more conveniently on
various detectors. The size and f# of the primary mirror strongly drive the instrument massand envelope. The primary mirror size also drives the design of the articulated mirror
mechanism and associated requirements. The optical design may be further optimized for the
wavelength ranges and detectors selected. The coatings of the mirror and gratings could also
be optimized for the particular application.

3-13

Appendix 3-A. EUVS count rates for the brightest lines
Count rates were calculated for the EUVS instrument using the CHIANTI spectral code forthe quiet Sun, active region and flare plasmas. The counts listed are for the total emission line
in a 0.33" spatial bin. The emission lines are sorted by channel and ionization temperature.
The ion and associated ionization temperature are included for reference. For the quiet Sun,
all lines with a total count rate of >10 counts/s are included. For active regions, emission
lines with a total count/s of >50 were tabulated. For flare plasmas, those emission lines withcount rates >10000 counts/s are included. Many valuable faint diagnostic lines have not beentabulated.
QUIET SUN INTENSITIES
Ion

Wvl.(A)

Tmax

Count/s

Fe IX
Fe X
Fe X
Fe XI

171.073174.531
177.240
180.408

5.90
6.10
6.10
6.10

86.9
36.4
16.4
14.1

He I
O V
Mg X

584.335
629.732
609.794

4.40
5.40
6.10

230.7
76.5
13.0

O III
N IV
O IV
O IV
Ne VIII
Ne VIII

703.854765.147787.710790.199770.410780.325

4.90
5.20
5.20
5.20
5.80
5.80

14.7
29.5
24.9
45.1
23.5
11.9

H I
H I
H I
C II
C II
C III
N III
O VI
O VI

949.745972.5381025.724
1036.339
1037.020
977.020991.5771031.914
1037.615

4.20
4.20
4.10
4.30
4.30
4.90
4.80
5.50
5.50

71.2
107.3
453.1
87.5
102.2
98.9
13.6
69.6
34.9

The lines counted in this table have all a count rate greater than
10.0000

3-14

ACTIVE REGION INTENSITIES
Ion

Wvl.(A)

Tmax

Count/s

Fe VIII
Fe VIII
Fe VIIIFe IX *Fe IX
Ni XI
Ni XIINi XIINi XIIINi XIIINi XIVNi XIV
Fe X
Fe X
Ni XV
Fe X
Fe XI
Fe XI
Fe X
Fe XI
Fe XI
Fe X
Fe XII
Fe XI
Fe XI
Fe XIIFe XIIFe XIIFe XIII

167.486
168.173168.544168.647
171.073148.374
152.154154.162
157.729
164.150
170.500
171.370
174.531175.263
176.741
177.240178.060
180.408180.441181.137
182.169
184.537186.887
188.232
188.299
192.394
193.509
195.119202.044

6.00
6.00
6.00
6.10
6.10
6.20
6.20
6.20
6.30
6.30
6.30
6.30
6.20
6.20
6.40
6.20
6.20
6.20
6.20
6.20
6.20
6.20
6.20
6.20
6.20
6.20
6.20
6.20
6.20

71.0
118.0
58.7
59.8
1671.6
56.3
73.4
50.4
125.3
90.8
72.2
108.2
1408.2
153.2
71.2
637.8
57.7
900.6
62.9
59.0
128.8
110.7
91.5
195.6
71.0
108.4
212.0
277.7
85.1

He I
O IV
O V
Si XII
Si XI
Mg X
Mg X

584.335
554.513629.732520.666
580.920609.794624.943

4.40
5.20
5.40
6.30
6.20
6.20
6.20

936.4
50.4
822.3
94.8
51.6
1010.1
613.9

O III
O II
O II
O V
N IV
S V
O IV
O IV
Mg IX
Ne VIII
Ne VIII

703.854718.506
718.568
760.446765.147786.470787.710790.199706.036770.410780.325

4.80
4.60
4.60
5.40
5.20
5.20
5.20
5.20
6.10
6.20
6.20

125.8
134.1
86.7
80.1
195.3
67.6
184.4
334.0
70.6
636.5
319.7

H I
H I
H I
C II
C II
S II

949.745972.5381025.724
1036.339
1037.020
1045.765

4.20
4.20
4.20
4.40
4.40
4.40

143.0
336.1
1441.1
95.0
103.3
54.8

3-15

C III
N III
N III
C II
O VI
O VI
Si VIIIFe XVIII

977.020989.799991.5771010.373
1031.914
1037.615
944.467974.860

4.70
4.80
4.80
4.50
5.50
5.50
6.10
6.70

2728.3
86.6
159.3
53.1
1258.7
631.0
57.7
68.3

The lines counted in this table have all a count rate greater than
50.0000

3-16

FLARE INTENSITIES
Ion

Wvl.(A)

Tmax

Count/sFe XXIIIFe XXII
Fe XXII
Fe XXIV

132.906
135.791
156.019
192.029

7.20
7.10
7.20
7.20

16711.910684.5
10742.9
41466.6

He I
O III
O IV
O IV
O V
Mg X
Fe XX
Fe XIX584.335
599.590554.513609.829629.732
609.794567.867592.236

4.50
4.90
5.20
5.20
5.30
6.80
7.10
7.00

22806.821647.016856.3
13597.1
123868.6
15972.914098.6
24262.3

O III
O III
O III
O III
O III
O III
O II
O II
N IV
O V
S V
O IV
O IV
O IV
Fe XX
Fe XXI

702.337702.838
702.896702.900
703.851703.854718.506
718.568
765.147760.446786.470787.710790.112790.199721.559786.162

4.90
4.90
4.90
4.90
4.90
4.90
4.70
4.70
5.10
5.30
5.20
5.20
5.20
5.20
7.10
7.10

15646.4
14272.3
12071.4
19085.9
19602.958808.9
17162.6
11081.7
61374.2
26905.2
26186.6
71908.2
13861.8128603.8108802.5
16011.5

H I
C III
N III
N III
C II
H I
H I
C II
C II
O VI
O VI
Fe XVIII

972.538
977.020989.799991.5771010.373
1025.724
1025.725
1036.339
1037.020
1031.914
1037.615
974.860

4.50
4.80
4.80
4.80
4.60
4.50
4.50
4.60
4.60
5.40
5.40
6.90

18466.2
834613.1
29670.652456.010053.1
72905.936616.4
24846.749448.2
67592.833837.3
116048.3

The lines counted in this table have all a count rate greater than
10000.0

3-17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

4-1

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

4-2

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

4-3

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.

4-5

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

4-6

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.

4-7

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

4-9

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

4-10

(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.

4-11

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

4-13

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.

4-25

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.

4-26

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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Chapter 5:Spacecraft and Mission Design

The spacecraft characteristics as the baseline are summarized in Table 5.1. Key itemsof the characteristics are described in following sections.
Table 5.1. Characteristics of the baseline spacecraft.
Spacecraft layout
WeightSize
Launch vehicle
Orbit
Power
Command & Telemetry

Attitude control

Science operations

Telescopes mounted on the optical bench, which is on the Bus
module. See Figure 7.1.
4930 kg (at liftoff), 2330 kg (dry weight)
3.2 m x 3.2 m x 7.4 m, excluding solar array paddles
JAXA H-IIA (type 202)
An inclined geo-synchronous orbit (Baseline)
A sun-synchronous polar orbit (optional)~3,500 W (maximum)
Uplink and housekeeping downlink: S-band
Downlink: X-band 16-QAM (16Mbps)
Ka-band QPSK (80Mbps, optional)
3-axis body pointing with high accuracy
Image stabilization system in some telescopes
A function to change the pointing around Z-axis for matching
spectrograph slit direction to the observing target.
Hinode operation scheme, with a final target selection around 10
to 15 minutes before the uplink in cases where this is essential forthe science.

Two candidates exist for the X-ray imaging telescope. Here we assume thephoton-counting grazing-incidence telescope, because it will require larger size andmass and higher telemetry rate.

5.1. Spacecraft Layout
The layout concept of the Solar-C Plan-B spacecraft is shown in Figure 5.1.
Technical heritage acquired with the development of the Hinode (Solar-B) spacecraft isused as much as possible, including the design concept of the spacecraft structural
layout. The three telescopes are mounted with mounting legs on the spacecraft OBU(Optical Bench Unit), which is a CFRP-made cylindrical structure providing stablecondition to the telescopes. The telescope portion of the UV-Visible-IR telescope(SUVIT) is equipped inside the OBU cylinder. Its focal plane instruments and the othertwo telescopes are mounted on the external wall of the OBU. The mounting legs are
made from low-thermal expansion materials, and provide 6-axis freedom kinematic
mount to the telescopes. The OBU is mounted on the spacecraft bus module with a truss
structure in order to isolate the OBU and telescopes from the influence of the thermal
deformation of the bus module.
It is important to minimize co-alignment changes among the three telescopes andthe attitude sensors against the launch load and the in-flight thermal environment. The
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three telescopes and the attitude sensors (UFSS and IRU) are mounted on the samestructural component, i.e., OBU, to achieve precise co-alignment among them.
The size and mass of the telescopes used for the initial design are listed in Table 5.2.
The values in this table are the initial guess based on the aperture size and focal length
currently under consideration for each telescope. The values contain significant
uncertainty and will be revised as the design is in progress.

Figure 5.1. SOLAR-C plan-B spacecraft
Table 5.2. Size and weight of the telescopes used as the initial parameters.
Instruments
Telescope
UV-Vis-IR
Spectrographtelescope
ImagerUV/EUV spectroscopic telescope
X-ray telescope (photon counting)
Total

Size (mm)ϕ2300 x 4300
1000 x 400 x 3200
1000 x 400 x 3200
1000 x 500 x 5000
400 x 400 x 4500

Weight (kg)
500
100
100
150
150
1000

The initial mass estimate result is shown in Table 5.3. The total mass of all thetelescopes (1,000 kg) is a very rough value (upper-side value) and this will be revised as
the telescope design is in progress. The spacecraft bus system is about 1,260 kg. The
propellant for moving the spacecraft from a geo-transfer orbit (GTO) into an inclined
geo-synchronous orbit (GEO) and maintaining the GEO orbit is counted as about 2,300
kg, which is the initial guess based on the information available on the website ofMichibiki (QZSS), which has an inclined geo-synchronous orbit (GEO). The reasons
why we are considering a GEO orbit as the primary option can be found later.
JAXA H-IIA 202 rocket is assumed as the baseline launch vehicle. This rocket iscapable of installing a 4-ton spacecraft into a GTO orbit. The initial mass estimateshows about 500 kg over the 4-ton, and we will need to make efforts to reduce the mass
by looking at the accuracy of mass estimates. Note that H-IIA 204 rocket, with foursolid rocket boosters, can install a 6-ton spacecraft into a GTO orbit, although it is much
5-2

more expensive.
Figure 5.2 shows the dimension of the spacecraft, which can be fitted in the
Model-4S fairing envelopes. The height is about 7.4m. The front-end position of the
X-ray and UV/EUV telescopes may be located much below the front position of the
UV-Visible-IR telescope, if the 4S fairing is used. This may give a contamination risk to
the UV/EUV and X-ray telescopes, because these telescopes entrance views the outerside surface of the UV-Visible-IR telescope. In the design phase, we should carefully
evaluate this layout and contamination issue in order to minimize the contamination
risk.
Table 5.3. Initial mass estimate result.
Components
Mission instruments
Structures
Bus module
Optical Bench
Thermal control system
Data handling (space wire) system
S and X-band
Communication
Ka-band (optional)
Electric power system
Solar array paddles

Weight (kg)
1000
340
70
100
10
21
TBD
142
120

Attitude & Orbit control system

100

Chemical propulsion system

170

Other peripheral hardware, such as
electrical cables and mechanical parts
Total (dry weight)
Propellant

Remarks
Table 7.2
Based on a standard Bus module
Sizing from Hinode optical bench
Based on Astro-H (space wire) system

Based on Section 7.2.3.

2 wings, 2 panels for each wing.
Including MIB (momentum wheel
isolation bench, 30kg)
Including apogee kick motor (~500N
thruster, section 7.3)

185
2258
2300

Total (at liftoff)

Based on the information available on theMichibiki (QZSS) website.

4558

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Figure 5.2. The spacecraft layout in the 4S fairing of the JAXA H-IIA rocket

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5.2. Science Telemetry
5.2.1. Data rate estimate
Table 5.4 summarizes the data rates estimated from the three mission instruments.
This is the initial guess based on the preliminary science design including the format
size of detectors and typical exposure cadence. Details of the estimate can be found in
Appendix 7-A. The table tells that high rate science telemetry is required to acquire
spectroscopic and polarimetric data with high cadence and resolution. “Standard” caseis for continuous observations without interruption, producing data with a constant rate.
“Burst” case means that the data is produced with highest data rate, but the duration ofsuch observations is limited.
The minimum requirement is that the data can be acquired constantly with the data
rate (the rate after compression) of about 8 Mbps in total. According to evaluation of
Hinode data, lossy compression is acceptable to most science cases, excepting X-ray
telescope’s photon counting data, although the effect from lossy compression needs tobe evaluated in the future. Thus, we assume to apply lossy compression to the data.
Also, much higher rate needs to be handled on board to have burst caseobservations. Assuming lossy compression can be used for burst case, the data output
with 570 Mbps (before compression) and 143 Mbps (after compression) should be
handled properly on board.
Table 5.4. Estimated data output rates from the mission instruments.
Instruments
Solar UV-Vis-IR telescope
UV/EUV spectroscopic
telescope
X-ray telescope(photon counting)Total

no comp.
14.4
320
7.0
140
6.9
110
28.3
570

standard
burst
standard
burst
standard
burst
standard
burst

data rate (Mbps)
lossless comp.
lossy comp.
6.1
2.5
128
48
3.0
1.5
60
30
3.5
55
12.6
7.5
243
143

5.2.2. Data handling on board and data recorder
The overall flow of the science data on board is described in Figure 5.3. It is thebaseline that data compression is performed in each telescope before producing the data
into the spacecraft bus. Since different types of data products (imaging, spectroscopic,
and photon counting) are generated in the three mission instruments, it would be good
to use different compression algorithms more suitable for the properties of data products.
After compressing and packetizing the data, each instrument sends compressed data to
the spacecraft. SpaceWire may be used for the data interface between the telescopes andspacecraft. There is no mission data processor, which is different from Hinode science5-5data handling system.
When we run continuous (i.e. 24 hours every day) observations with the 8Mbps
data generation rate, the total data amount is about 700 Gbits per day. Since themaximum duration of the ground station support at USC may be less than 8 hours, thedata needs to be recorded in the data recorder. The capacity of the data recorder should
be enough to allow us to have continuous observations for 24 hours every day.
Moreover, the data recorder allows us to run the burst observations; Data produced
burstly is stored in the data recorder, followed by downloading the stored data after theacquisition when the spacecraft is linked to a ground station. The capacity of the data
recorder will be determined by considering the observation plans and the capacity
reality.

Figure 5.3. Overall data flow estimate

5.2.3. Science Telemetry System
The high-speed telemetry channel is absolutely required for the Solar-C Plan B
mission to transfer the huge amount of data expected from the onboard telescopes. The
requirement is to design the telemetry and data storage system to enable downloading
the volume of the data produced constantly with the rate of about 8 Mbps.
An initial trade-off study suggests either an X-band telemetry system (16Mbps,
16-QAM modulation, with semi-directional low-gain antennas), or Ka-band telemetry
system (80[TBD]Mbps, QPSK modulation, with directional high-gain parabolaantenna), or the both of them. Each telemetry system is briefly described below.

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X-band telemetry system (Figure 5.4)
X-band telemetry is easy to handle, but the problem is that the allowable bandwidthis restricted within 10MHz for near-earth space science applications in X-band
frequency band, according to a recommendation by the ITU’s SFCG (Space Frequency
Coordination Group). If we use the QPSK (phase shift keying) modulation, same as
Hinode’s X-band modulator, the telemetry speed needs to be below about 8 Mbps tomeet the signal bandwidth restriction. Instead, the 16-QAM (quadrature amplitude
modulation), which is widely used on the ground communication and will be tested in
one of Japanese technology demonstration satellites, can increase the telemetry speed to
16 Mbps.
If we have 12 hours for the telemetry downlink in total per day, the telemetry system
can transfer all the data acquired with 8Mbps continuously for 24 hours. We assume
USC 34m antenna as the primary antenna, and the system can be designed to have apositive link margin. Considering other spacecraft operations at USC 34m, the durationof downlink may be restricted below 8 hours. Thus it would be helpful to haveadditional downlinks on NASA and ESA stations, if antennas are available in Asia,
Oceania, and Pacific Ocean area. It is also noted that a huge mass memory (recorder) isneeded to record the data taken during non-contact period.

Figure 5.4. X-band telemetry system
Ka-band telemetry system
Another possibility is to newly explore the usage of Ka-band in Japan. Higher
telemetry rate is possible because a wider bandwidth can be utilized in the Ka band by aspacecraft. Two candidate bands are identified; 1) ~17GHz, a portion of groundequipments available at USC can be used to minimize the cost. 2) ~26GHz, this band isused at the Near Earth Network (NEN) of NASA. It is needed to newly prepare thesystem for receiving signals at this band in Japan, if the telemetry is received at Japan.
When the bandwidth is 100 MHz, 80 Mbps is achieved with QPSK. The downlink
5-7

duration can be reduced to a few hours, although a more huge mass memory (recorder)
is needed.
A concern is that attenuation of signals due to bad weather (rains and heavy clouds)
in Japan affects quality of the data and gives complication to real-time daily operations.
The both X-band and Ka-band systems work in case of the GEO orbit; X-band
16-QAM (16Mbps) as a baseline and Ka-band QPSK (80Mbps) as an option, though wehave to implement the Ka-band capability at a ground station. However, it is impossible
to adapt the X-band system if the spacecraft is in a sun-synchronous polar orbit, because
about 15 passes are too small to download all the data. The only way is to adapt theKa-band system, but a Ka-band antenna is required in the polar region, such as at
Svalbard. At least 144 minutes are needed as the total contact duration to download all
the data with 80Mbps and this duration corresponds to about 12 station contacts. During
each contact (duration ~ 12 minutes), the Ka-band high-gain antenna is quickly moved
to direct to the ground station, resulting in a concern of micro-vibration from theantenna mount mechanism. More details of the downlink system for the GEO orbit aresummarized below:
1.

2.

GEO/X-band (baseline) The 16-QAM modulator, which is under development for the ASNARO
mission The bit rate is 16 Mbps, which occupies the band width ~10MHz
 Semi-directional low-gain antenna (angle ~ 60°, 0dbi), which is used in
HAYABUSA. At least 6 (TBD) antennas are mounted to cover any direction,
one of which is selected by switch.
 50W TWTA amplifier, which is under development for the ASTRO-G mission.
 The link margin is XXX dB for USC 34m.
 It is necessary to implement the X-band 16-QAM demodulator at the USC.
GEO/Ka-band (optional) The QPSK modulator and the bit rate is 80 Mbps (band width=100MHz)
 ϕ 0.3m high-gain antenna (32dbi) + 0.5W SSPA, whose pointing is driven by a
two-axis mount mechanism.
 The link margin is XXX dB for USC 34m.
 It is necessary to make the Ka-band (17GHz) at the USC available. Anotherpossibility is to make use of Ka-band (26GHz) with NASA or ESA groundstation.

5.3. Orbit
As briefly discussed in section 5.2.3, two orbit candidates are considered for thePlan B spacecraft. Currently, an inclined geo-synchronous orbit (GEO) is considered as
the baseline, with a sun-synchronous polar orbit (SSO) as the backup.
Geo-synchronous orbit (GEO): Same as Solar Dynamic Observatory (NASA)
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 Altitude: 36000km, inclination: <30°, period: 1 day
Sun-synchronous polar orbit (SSO):
 Same as Hinode, TRACE, and IRIS.
 Altitude: 680-800 km, inclination: 97-98°, period: 98 min
The GEO orbit would provide much more advantages than the SSO orbit,
although a demerit on the weight of the spacecraft. Advantages are listed below:
 X-band telemetry can be utilized to download the data volume required forobservations, though the duration of contacts is fairly long (8 hours or longer perday).
 Thermal environment is colder in GEO than in SSO because of lower influence ofinfrared radiation from the earth, which much helps the thermal design of the
optical telescope (see section 4.3.4).
 GEO is thermally very stable, providing less orbital variation of thermal
environment. This would reduce structural deformation, resulting in less orbital
variation in pointing stability (section 5.4.1) and temperature fluctuation oftelescope optics.
 A new type of methodology for science operations (section 5.7) can be adapted,
which would much help the narrow field-of-view spectroscopic telescopes toperform optimized observations
 Eclipse seasons are shorter in GEO than in SSO. SSO has eclipse season with
about 20 mins (at maximum) duration every 98 mins orbital period and the season
continues three months every year. The GEO’s eclipse is about 72 mins (atmaximum) duration per day and it continues only one month, twice every year (seeFigure 5.5). A potential demerit is to need a large capacity of the battery forkeeping some components warm without solar heat input.

Figure 5.5. Eclipse duration in the case of GEO. The maximum duration is 72 mins perday. The eclipse season happens twice every year around the equinox. The duration of

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the eclipse season depends on the inclination angle of the orbit.

Figure 5.5: Visibility of the spacecraft from USC that is located at latitude of 30°N
when the spacecraft is in GEO with inclinations.
It is not necessary to install the spacecraft into the geostationary orbit (GSO) in which
the spacecraft is located directly above the equator where the inclination angle is 0°.
When the orbit is at non-zero inclination, the duration of the eclipse season can be madeshorter than one month (Figure 5.5). The spacecraft is no longer stationary when viewed
from a ground station at the non-zero inclination. The spacecraft is always visible at anelevation angle higher than 45° at USC (latitude 31°N) when the inclination angle isless than 30° (Figure 5.6), which is high enough for continuous contact without
interruption.
The most critical issue of GEO is that an apogee kick motor is required to install the
spacecraft from GTO to GEO. The engine and its propellant take half of the total weightat the lift-off as is shown in Table 5.3.

5.4. Pointing Stability and Attitude Control
5.4.1. Requirements on the pointing stability
We set targets of angular resolution better than 0.1 arcsec in the Solar
Ultraviolet-Visible-IR telescope, and better than 0.2 - 0.5 arcsec in the UV/EUV and
X-ray telescopes. The angular resolution requires high stability of the body pointing,otherwise images are blurred due to residual pointing jitter. The requirements on the

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pointing stability are thus derived based on the requirements on the angular resolution
and exposure duration to obtain a data set (an image or set of images), and aresummarized in Table 5.5. In Figure 5.7, the requirements on the pointing stability forboth HINODE and SOLAR-C are indicated and compared as a function of frequency.
Table 5.5. Requirements on the pointing stability of the mission instruments.
Instruments

Time scale

Solar UV-Vis-IR telescope (*)

1 sec
10 sec1 hourMission life
0.5 sec
5 sec
1 hourMission life
1 sec
10 sec1 sec
1 min
1 hourMission life

UV/EUV spectroscopic
telescope

X-ray
telescope

Normal incidence(*)

Grazing incidence

Instruments

Time scale

Requirements(θx/θy)
0.015
0.015
2
20
0.3
0.3
2
32
0.1
0.1
0.3
0.7
8
32
Requirements
(θz)
50

Unit
arcsec 3σ
arcsec 3σ
arcsec 0-p
arcsec 0-p
arcsec 3σ
arcsec 3σ
arcsec 0-p
arcsec 0-p
arcsec 3σ
arcsec 3σ
arcsec 3σ
arcsec 3σ
arcsec 0-p
arcsec 0-p

Unit

Solar UV-Vis-IR Telescope
1 hourarcsec 0-p
UV/EUV spectroscopic
1 hour100
arcsec 0-p
telescope
Normal incidence
1 hour50
arcsec 0-p
X-ray
telescope
Grazing incidence
1 hour100
arcsec 0-p(*) The pointing stabilities are to be achieved with the usage of an image stabilization system inside
the telescopes.

It is important to identify critical items to achieve the pointing stability
requirements by comparing on-orbit pointing performance of HINODE that achieved
the highest pointing stability in Japanese space missions. We evaluate the pointing
stability of HINODE in the frequency domain using telemetry data taken with the Ultra
Fine Sun Sensor (UFSS) and the mechanical gyroscope (Inertia Reference Unit, IRU)
that are components of the attitude and orbit control system (AOCS) on-board HINODE.
They provide residual pointing errors with a temporal sampling of 0.5 sec, which
provides the pointing stability in the frequency domain lower than 1 Hz. The correlation
tracker (CT) on-board HINODE provides signals of pointing errors with a temporal
sampling of 580Hz, which is used to evaluate high frequency disturbances in thefrequency range between 0.1 Hz and 290 Hz. The on-orbit pointing stability evaluatedwith UFSS, IRU, and CT of HINODE is shown in Figure 5.7, as comparison with therequirements on the pointing stability for HINODE (top) and SOLAR-C (bottom). It isfound that the AOCS on-board HINODE delivers better performance than therequirements in the frequency range lower than 0.1 Hz (upper limit of the AOCS control
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band). It is important to note that, without the image stabilization system, it is difficult
to realize the good spatial resolution with SOT because the pointing stability of about
0.1 arcsec (0-p) at 0.1 Hz achieved with the AOCS is worse than the requirements of
SOT. The image stabilization system has a crucial role in the reduction of the attitude
jitter in the frequency range between 10-2 Hz and 10 Hz.

Figure 5.6. On-orbit performance of the pointing stability in HINODE in comparison
with the requirements of the pointing stability for the mission instruments aboard
HINODE (top) and SOLAR-C (bottom). The thick, dashed and dotted lines represent
the requirements as a function of frequencies while the lower boundaries of the shaded
belts indicate conditionally acceptable level, which is defined as one third of therequirements. The blue filled circles show the pointing stability evaluated with UFSS,the red open circles are that with IRU. The open cyan squares indicate the pointing
stability evaluated with CT in the case of servo-on, and filled orange squares are that in
the case of servo-off.

By comparing the on-orbit pointing performance of HINODE with therequirements for SOLAR-C, we have identified three critical items for further
improvements necessary in the Plan B spacecraft:

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(1) Orbital variation (10-3 – 10-5 Hz)
It is caused by misalignment between the attitude sensors and the telescopes due tothermal deformation of structures supporting them. In the case of SSO, temperaturefluctuation is mainly induced by significant influence of infrared radiation from theearth. In SOLAR-C, the orbital variation is to be reduced significantly thanks to thelower temperature fluctuation at GEO.
(2) Degradation of the pointing stability at around 0.1 Hz
There is an unexpected degradation at around 0.1Hz in the on-orbit pointing
performance of HINODE. The frequency of 0.1 Hz corresponds to the upper limit ofcontrollable frequency of the AOCS. The stability is 0.1 arcsec (0-p), which is notenough to achieve the spatial resolution better than about 0.3 arcsec. It is necessary
to clarify the cause and to improve the AOCS not to induce the degradation in
SOLAR-C. If it is difficult to improve the pointing stability at around 0.1Hz onlywith the AOCS, the mission instruments have to be equipped with an image
stabilization system inside them.
(3) High-frequency disturbances (>100 Hz)
The disturbances at the frequencies higher than 100 Hz have to be suppressed toachieve the resolution of 0.1 arcsec in SUVIT of SOLAR-C. They are induced bystructural resonance between the telescope structures and disturbances of momentum
wheels (MWs) and mechanical gyros. Because the image stabilization system aboard
HINODE is not effective to such high-frequency (>20 Hz) disturbances, one possible
approach is to improve the image stabilization system to have a broader response in
the frequency as high as 100-200 Hz. The other approach is to suppress generation
or transmission of the disturbances from MWs and the mechanical gyroscopes,
which is technically possible if MWs are mounted on a vibration isolation bench.

5.4.2. Attitude control system
The attitude control system provides three-axis stabilization of the spacecraft to point
to an observing target on the solar surface. We can basically use the same configuration
as HINODE, which achieved the highest pointing stability in Japanese space missions,but some modification is necessary in order to improve the items required for further
improvements of the pointing stability described in section 5.4.1. Major components inthe attitude control system are listed below.


Attitude and orbital control processor (AOCP)
 The attitude computer is based on Hinode and other JAXA spacecrafts.
Ultra Fine Sun Sensor (UFSS) The sensor precisely measures the direction of the Sun with sub arc second
accuracy.
 The performance of UFSS on-board HINDOE is enough for the Solar-C
requirements.

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









Fine Sun Sensor (FSS) The planned sensor has been used on Planet-C and so on.
Star Tracker (STT) The start tracker measures the roll around the direction of the Sun.
 The performance of the star trackers used in other spacecrafts meets the
Solar-C requirements.
Tuned Dry Gyroscope (TDG) The gyroscope measures the angle change of the spacecraft attitude as aninertia reference unit (IRU).
 TDG has been used in many space missions in Japan. We have to investigate
the possibility to reduce the disturbances caused by the mechanical gyroscope.
 As an alternate option, fiber optic gyroscope (FOG) is considered for a
disturbance-free gyroscope.
Geomagnetic Aspect Sensor (GAS) The sensor may be purchased from an aerospace company.
Momentum wheels (MW) Momentum wheels are used as actuators of the spacecraft attitude.
 The baseline for momentum wheels is same as used in Hinode.
 For suppressing transmission of the micro-vibration disturbances, vibration
isolation bench may be required to mount momentum wheels.
Magnetic torquer (MTQ) The sensor may be purchased from an aerospace company (following theprevious missions).

5.4.3. Flexibility on the role attitude
In the Hinode case, the spacecraft role is always fixed to align the spacecraft Y-axis,
i.e., Y direction of CCD detectors, to the solar North-South direction. For the Solar-C
Plan B mission, it would be useful to allow changing the spacecraft role attitude,
depending on the nature of observing target. For example, when the slit of spectroscopic
telescopes is placed along a coronal loop structure, the spectroscopic telescopes canrecord spectral lines as a function of spatial position along the loop structure with very
limited number of slit scan positions (or in sit-and-stare mode). This can provide highcadence of observations, allowing us to track energy flows along magnetic fields.
This capability will be considered in the AOCS functional design. It would be notdifficult in implementing this capability from viewpoint of AOCS. However, thiscapability would give complication to the command uplink and telemetry downlink
operations. We may need more number of onboard antennas to cover any directions forthe communication link with a ground station. Or, insufficient communication link levelmay be given depending on the roll angle.
5.5. Spacecraft thermal design to dump heat load from 1.5m telescope
In the thermal design of the spacecraft system, we basically follow the same concept
as adapted in HINODE, which is summarized in the following:
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



The OBU is made of low-thermal expansion and high-thermally conductive CFRP.
The telescopes are thermally isolated from OBU and the spacecraft bus.
The OBU is thermally isolated from the spacecraft bus.

However, as described in section 4.3.4, some part of the large heat load given to
the primary mirror (M1) of the UV-Visible-IR telescope (SUVIT) needs to betransferred toward the spacecraft bus direction (-Z direction). With having the heat path
toward the bus direction, the temperature of M1 can be kept well within the acceptabletemperature range. Thus, a thermal interface is required between the SUVIT’s primary
mirror and the spacecraft. A preliminary thermal concept is to have heat pipes to transfer
the heat to the radiators equipped at the bus side (Figure 5.8). The heat load at M1 isdumped to a cold plate located behind M1 with radiation. Heat pipes connect the coldplate to the radiators on the side of the bus or OBU structure. The heat load of about 100– 200 W is required to transfer via heat pipes. A preliminary study shows that the size ofthe radiators necessary for 100-200W heat dump can be accommodated either on theside of the bus module or OBU structure.

Figure 5.7. Thermal design concept to dump the heat load from M1 of the SUVIT-OTA.
5.6. Spacecraft power system
The spacecraft has solar array panels that generate the electric power. Considering
the size of the Bus module and telescopes, 4 solar array panels would be a baseline.
They are attached on the Bus; either two wings, each of which has two panels, as shownin Figure 5.1, or 4 panels (wings) at each face of the Bus module.
The 4 panels can generate the power up to ~3500W. The power budget will be
evaluated in the design phase, but we will have no critical concerns on the power budget,from our former experience on Hinode (~1100W).
The proposed orbit provides a long continuous duration of non-eclipse condition.
For eclipse periods and initial launch operation, the power system needs a power battery
to provide the electrical power to the onboard components. The size of buttery will be
determined after the initial evaluation of power budget in the design phase.

5-15

The power system provides an unregulated 50V power to all the components onboard; 50-52 V in non-eclipse period and 30-49V in the eclipse period, which is same asused in Hinode.
5.7. Science operations
Since the spacecraft is viewed for a long duration from a ground station (USC),
science daily operations may become more flexible than Hinode’s operation. The
following methodology may be added to the science planning scheme; an observer inthe operation room checks images in real time, and make commands for adjusting thetelescope pointing to the small area of interest, and then the operator uplinks thecommands to the spacecraft during a contact (Figure 5.9). This kind of operations would
be very powerful, because telescopes may have the field of view smaller than Hinodeand the slit position for spectroscopic telescopes is very critical to capture thephenomena of interest.

Figure 5.8. A new methodology of science operations that can be considered in the
GEO case
5-16

Appendix 5-A. Data rate estimate
We here estimate data production rates from the mission instruments using their preliminary
design including formats of their detectors and typical exposure duration in order to clarify how
much telemetry rate is required to download the data from the spacecraft to the ground. We define
two typical observing cases for each instrument; one is for continuous observations without
interruption (“standard case”), which produces a lower data rate, and the other is for short-term burst
observations (“burst case”), which produces a higher data rate. In the estimation, we assume a bit
depth of their data products, and how much the data amount can be shrunk with lossless and lossy
compression algorithms. Details are summarized in the following.


Solar UV-Vis-IR telescope (SUVIT) / Spectrograph package
 Standard case:- In this case, we can achieve high signal-to-noise ratio of about 104 with deep
integration for measuring chromospheric magnetic fields.
- Number of pixels of an image set:1024 (spatial) x 512 (wavelength) x 4 (Stokes) x 2 (orthogonal pol.)= 4 Mpixel- Bit depth:
16bit/pixel (no comp.), 7bit/pix (lossless comp.), 3bit/pix (lossy comp.)- Temporal cadence:
10 sec at each slit position
 6.4 Mbps (no comp.), 2.8 Mbps (lossless comp.), 1.2 Mbps (lossy comp.)
 Burst case:- In this case, we can achieve standard signal-to-noise ratio of about 2x103 with 1 sec
integration for rapid slit scanning.
- Number of pixels of an image set:
2048 (spatial) x 1024 (wavelength) x 4 (Stokes) x 2 (orthogonal pol.)= 8Mpixel- Bit depth:
16bit/pixel (no comp.), 6bit/pix (lossless comp.), 2bit/pix (lossy comp.)- Temporal cadence:
1 sec at each slit position
 128.0 Mbps (no comp.), 48.0 Mbps (lossless comp.), 16.0 Mbps (lossy comp.)

Solar UV-Vis-IR telescope (SUVIT) / Imager package
 Standard case:- In this case, we get minimum wavelength sampling for Dopplergrams with thenarrow-band filtergraph with reduced spatial and temporal sampling.
- Number of pixels of an image set:
2048 x 2048 (spatial) x 5 (wavelength)= 20 Mpixel

5-17

- Bit depth:
12bit/pixel (no comp.), 5bit/pix (lossless comp.), 2bit/pix (lossy comp.)- Temporal cadence:
30 sec for each wavelength scan
 8.0 Mbps (no comp.), 3.3 Mbps (lossless comp.), 1.3 Mbps (lossy comp.)
 Burst case:- In this case, we get full wavelength sampling of a spectrum line for
imaging-spectroscopic observations with the narrow-band filtergraph with full spatial
and temporal sampling.
- Number of pixels of an image set:
4096 x 4096 (spatial) x 10 (wavelength)= 160 Mpixel- Bit depth:
12bit/pixel (no comp.), 5bit/pix (lossless comp.), 2bit/pix (lossy comp.)- Temporal cadence:
10 sec for each wavelength scan
 192.0 Mbps (no comp.), 80.0 Mbps (lossless comp.), 32.0 Mbps (lossy comp.)


UV/EUV high-throughput spectroscopic telescope Standard case:- In this case, we get a reduced set of spectrum lines with deep integration.
- Number of pixels of an image set:
2048 (spatial) x 128 x 10 (wavelength)= 2.5 Mpixel- Bit depth:
14bit/pixel (no comp.), 6bit/pix (lossless comp.), 3bit/pix (lossy comp.)- Temporal cadence:
5 sec at each slit position
 7.0 Mbps (no comp.), 3.0 Mbps (lossless comp.), 1.5 Mbps (lossy comp.)
 Burst case:- In this case, we get a full set of spectrum lines with shorter integration for rapid slit
scanning.
- Number of pixels of an image set:
2048 (spatial) x 256 x 20 (wavelength)= 10 Mpixel- Bit depth:
14bit/pixel (no comp.), 6bit/pix (lossless comp.), 3bit/pix (lossy comp.)- Temporal cadence:
1 sec at each slit position
 140.0 Mbps (no comp.), 60.0 Mbps (lossless comp.), 30.0 Mbps (lossy comp.)

5-18



X-ray imaging (spectroscopic) telescope Assumptions:- Perform tagging for each photon-incident event- Tagging information
 Pixel address (X, Y) Photon energy Incident time
- No on-chip/on-board pixel summing. Perform on ground if necessary.
- Data compression may be possible, which reduces data rate to half.
 Standard case:- In this case, we get photon counting imaging spectroscopy in 128 x 128 pixels ROI.
- Pixel address (X, Y):
7 bits x 2 = 14 bits (for 0–127)- Photon energy:
10 bits (0–1023 ch)- Exposure time:
every 1 ms
- Ratio of event pixels:
2 % (i.e., 128^2 x 10% = 327.7 pixels contain events)
 6.9 Mbps (no comp.), 3.5 Mbps (comp.)
 Burst case:- In this case, we get photon counting imaging spectroscopy in 512 x 512 pixels ROI.
- Pixel address (X, Y):
9 bits x 2 = 18 bits (for 0–511)- Photon energy:
10 bits (0–1023 ch)- Exposure time:
every 1 ms
- Ratio of event pixels:
10 % (i.e., 512^2 x 10% = 6553.6 pixels contain events)
 110.0 Mbps (no comp.), 55.0 Mbps (comp.)

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R-9

Acronyms
16-QAM

16 Quadrature Amplitude Modulation

2FS

Secondary Field Stop

3D

Three dimensional

A/D

Analog to Digital

AIA

Atmospheric Imaging Assembly, SDO

ANNartificial neural network

AOCS

Attitude and Orbit Control System

APS

Active Pixel Sensor

AR

Active region

ASNARO

Advanced Satellite with New System Architecture for Observation

BF

Broadband filtergraph

BOL

Beginning of life

BS

Beam splitter

CCD

Charge Coupled Device

CDS

Coronal Diagnostic Spectrometer, SOHO

CFRP

Carbon Fiber Reinforced Plastics

CH

Coronal hole

CHIANTI

An Atomic Database for Spectroscopic Diagnostics of Astrophysical Plasmas

CME

Coronal Mass Ejection

CMOS

Complementary Metal Oxide Semiconductor

CMU

Collimator Mirror Unit

CRD

complete angle and frequency redistribution

CT

Correlation Tracker

CTE

Coefficient of thermal expansion

DEM

Differential Emission Measure

EIS

Extreme ultraviolet Imaging Spectrometer, HinodeEIT

EUV Imaging Telescope, SoHOEOLEnd of life

EQ

Equilibrium
R-10

ESA

European Space Agency

EUNIS

Extreme Ultraviolet Normal Incidence Spectrograph

EUV

Extreme Ultraviolet

EUVS

EUV/FUV high-throughput spectrometer

FOV

Field of View

FPGA

Field-Programmable Gate ArrayFUVFar Ultraviolet

FWHM

Full width half maximum

GEO, GSO

Geo-Synchronous Orbit

GOES

Geostationary Operational Environment Satellite

GTO

Geo-Transfer Orbit

HAZEL

Hanle and Zeeman Light (Inversion code)

HDM

Heat Dump Mirror

HPD

Half power diameter

IR

Infrared

IRU

Inertial Reference Unit

IRISInterface Region Imaging Spectrograph

ISS

Image stabilization system

ITU

International Telecommunication Union

JAXA

Japan Aerospace Exploration AgencykGkilo Gauss

LDE

Long-duration flare event

LVCR

Liquid crystal variable retarder

M1

Primary mirror

M2

Secondary mirror

MET

Microfield Exposure Tool

MHD

Magneto-Hydro-Dynamics

MIB

Momentum wheel Isolation BenchMK

Million Kelvin

MW

Momentum Wheel

NASA

National Aeronautics and Space Administration

R-11

NEN

Near Earth Network

NF

Narrowband filtergraph

NGXT

Next generation X-ray telescope

NI

Normal incident

NIR

Near infrared

NLFFF

non-linear force free field

OBU

Optical Bench Unit

OTA

Optical Telescope AssemblyPC

Photon counting

PCA

Principal component analysis

PCU

Polarization Calibration Unit

QE

Quantum efficiency

QPSK

Quadrature Phase Shift Keying

QS

Quiet Sun

RAISE

Rapid Acquisition Imaging Spectrograph Experiment

RHESSI

Reuven Ramaty High Energy Solar Spectroscopic Imager

RMS

Root mean square

SDO

Solar Dynamics Observatory

SFCG

Space Frequency Coordination Group

SIR

Stokes inversion based on response function

SMEX

Small Explorer

S/N

Signal-to-Noise

SOHO

Solar and Heliospheric Observatory

SOT

Solar Optical Telescope, HinodeSP

Spectro-Polarimeter

SSO

Sun-synchronous polar orbit

SSPA

Solid State Power Amplifier

STEREO

Solar TErrestrial RElations Observatory

SUMER

Solar Ultraviolet Measurements of Emitted Radiation, SOHOSUVIT

Solar UV-Visible-IR Telescope

SXT

Yohkoh Soft X-ray Telescope

R-12

TBD

To be determined

TF

Tunable filter

TRACE

The Transition Region and Coronal Explorer

LTE

Local thermodynamic equilibriumTM

Tip-tilt mirror

TRTransition Region

TRACE

Transition Region and Coronal Explorer

TWTA

Travelling Wave Tube Amplifier

UFSS

Ultra Fine Sun Sensor

USC

Uchinoura Space Center

UV

Ultraviolet

UVCS

Ultraviolet Coronagraph Spectrometer, SOHOVERISVEry high angular Resolution Imaging Spectrometer

WFS

Wavefront sensor

WP

Wave plate

XIT

X-ray Imaging Telescope

XRT

X-ray Telescope, HinodeR-13