SOLAR-C Interim Report:

SOLAR-C Plan-A
A Direct Approach to the Origin of the Solar Activity Cycle

ISAS/JAXA SOLAR-C Working Group

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Contents
Executive Summary .................................................. 5
1. SOLAR-C Mission Science Goals ...................................... 10
1.1 Origin of Solar Magnetism .............................................. 10
1.1.1 Solar Magnetic Cycle ................................................. 10
1.1.2 Generation and Emergence of Magnetic Field ............................... 11
1.1.3 Dynamics of the Solar Convection Zone and Tachocline ........................ 14
1.1.4 Probing the Solar Interior ............................................. 15
1.1.5 Prediction of Solar Cycles ............................................. 17
1.1.6 Scientific Questions to be Solved by SOLAR-C ............................... 18
1.2 Exploration of Solar Interior and Solar Magnetic Activity ........................ 20
1.2.1 Q1: How is the global, cyclic, solar magnetic field generated? ..................... 20
1.2.2 Q2: What is the nature of the flows in the polar regions of the Sun and how do they interact
with magnetism? ....................................................... 21
1.2.3 Q3: How does the radiative energy output of the Sun depend on latitude? ............ 23
1.2.4 Q4: How does magnetic activity shape the structure and evolution of the polar solar coronaand how does it affect the Earth? ............................................ 24

2. Scientific Requirements ............................................. 32
2.1 Scientific Requirements ................................................ 32
2.1.1 Requirements for Helioseismology Observables .............................. 32
2.1.2 Requirements for Magnetic Fields Observables .............................. 34
2.1.3 Requirements for TSI Observations ...................................... 35
2.1.4 Requirements for EUV imaging and Spectroscopy Observables ................... 35
2.1.5 Observables of inner heliospheric imaging (Optional) .......................... 36
2.1.6 Observables of in-situ measurements (Optional) .............................. 36
2.1.7 Summary of Required Data ............................................ 37
2.2 Science Payload ...................................................... 38
2.2.1 Helioseismic and Activity Imager ........................................ 38
2.2.2 EUV & X-ray Spectroscopic Imaging Telescopes ............................. 39
2.2.3 Solar Irradiance Monitor .............................................. 41
2.2.4 Inner Heliospheric Imager (Optional) ..................................... 41
2.2.5 In-Situ Instruments (Optional) .......................................... 42
2.2.6 Other Instruments (Optional) ........................................... 42
2.3 Spacecraft Requirements given from the Science Objectives ....................... 43
2.3.1 Attitude Control Requirements .......................................... 433

2.3.2 Time management ................................................... 44
2.3.3 Data Compression and Data Storage ...................................... 44

3 Spacecraft System ................................................. 46
3.1 Spacecraft System Requirements .......................................... 46
3.2 Orbit and Mission Profile ............................................... 46
3.2.1 Trajectory Options for SOLAR-C Plan-A .................................. 47
3.2.2 Orbit and Mission Profile of SEP Option ................................... 48
3.3 Spacecraft System Design ............................................... 52
3.3.1 Spacecraft Configuration .............................................. 52
3.3.2 Mass Budget ....................................................... 53
3.3.3 Power Budget ...................................................... 54
3.3.4 Communication Link Budget ........................................... 55
3.3.5 System Thermal Design ............................................... 55
3.4 Spacecraft Subsystems ................................................. 56
3.4.1 Data Handling System ................................................ 56
3.4.2 Communication System ............................................... 56
3.4.3 Power Supply System ................................................ 57
3.4.4 Attitude Control System .............................................. 59
3.4.5 Chemical Propulsion System ........................................... 59
3.4.6 Ion Engine System ................................................... 59
3.4.7 Structure System.................................................... 61
3.4.8 Thermal Control System .............................................. 62

Appendix ........................................................ 63
Appendix A:........................................................... 63
A.1 Exploration of α-effect and turbulent diffusion ................................ 63
A.2 Flows associated with flux emergence ...................................... 63
A.3 The inclination requirement ............................................. 65
Appendix B: Jupiter Option ................................................ 68
B.1 Orbit Trajectory ..................................................... 68
B.2 Mission Profile ...................................................... 68
B.3 Mass and Power Budget ................................................ 69
Appendix C: References .................................................. 71
Appendix D: Acronyms ................................................... 72

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Executive Summary
Scientific Background
Understanding the origins of solar magnetic activity has been at the forefront of solar and stellar
physics since the discovery of the 11-year sunspot cycle nearly two centuries ago. Unraveling thismystery has broad implications not only for promoting a deeper knowledge of the Sun itself but
also for understanding the Sun’s influences on the heliosphere, the geospace environment, andpotentially the Earth’s climate system. Such influences regulate space weather, with increasing
economic impacts on our technological society as our reliance on telecommunications systems,
power grids, and airline travel continues to grow. As a readily observable example of anastrophysical magneto-hydrodynamic (MHD) dynamo, the Sun also provides unique insights into
the generation of magnetic fields by turbulent plasma flows throughout the universe, from planetaryand stellar interiors to stellar and galactic accretion disks to interstellar clouds.
The global magnetic polarity of the Sun reverses during each 11-year sunspot cycle so that theoverall period of the solar magnetic activity cycle is 22 years. It is a formidable challenge tounderstand how such remarkable regularity arises from the highly turbulent conditions of the solar
convection zone and how magnetic flux emerges from the solar interior to energize the solar
atmosphere and power solar variability. Large-scale flows (differential rotation and meridional
circulations) established by turbulent convection, plasma instabilities, and nonlinear feedbacks allplay an important role, spanning many orders of magnitude in spatial and temporal scales.
Modern solar observations coupled with sophisticated theoretical and numerical models haveyielded important insights into many aspects of solar magnetism but the basic physical mechanisms
responsible for generating these fields are still not understood. To make great scientific progress in
our understanding the Sun and the fundamental problems of cosmic magnetism, there is no doubtthat we need to continue both theoretical and observational efforts. On the observational side,
measuring solar internal flows is of the greatest importance. For this task helioseismology has
proven to be a powerful tool.
Helioseismic measurements are based on surface wavefield data, normally and preferably temporalseries of photospheric Dopplergrams, which are then analyzed to probe the solar interior structureand flows. With the so-called global methods, the wavefield data are used to measure (mainly
acoustic) eigenfrequencies of the Sun. The eigenfrequencies are then analyzed, often by way of
inverse methods, to probe the solar interior for thermal and dynamical structure of high degrees ofsymmetry, such as the spherically symmetric distribution of sound speed, or differential rotation as
the axisymmetric component of flows. With new local methods, wavefield data are used to measure
local resonant properties or wave propagation time for a given pair of points, by cross-correlating
local wavefields. These travel-time data are then analyzed to probe the interior for local and/or
asymmetric structures, such as meridional flow, convection and flows around active regions.
Differential rotation and meridional flows have already been measured by such helioseismology
techniques, up to about 60 latitude with a typical uncertainty of the order of a m/s. It is essential toextend this measurement to the polar region, partly because without such measurement we will
never be confident of our understanding of dynamics of the Sun as a whole, and partly because thepolar region is where the magnetic flux reverses and the meridional flow, which plays an important5

role in carrying the magnetic flux, should turn in towards the solar interior, and where polarity
reversals take place in the surface layers.
Another related mystery is the total solar irradiance variation over the solar cycle. The total
irradiance of the other solar-like stars that exhibits activity at a level that is similar to the Sun, on
average varies around 0.3 per cent over their activity cycles. On the other hand, the solar irradiancevaries only by 0.1 per cent. There is a well-founded suspicion that the solar irradiance depends onlatitude, thereby creating a great interest in measuring the solar irradiance from high heliographic
latitudes.
Science Objectives
SOLAR-C Plan A, from its highly inclined orbit around the Sun, aims to combine helioseismic andmagnetic observations, solar irradiance measurements at various latitudes, and finally EUV and Xray observations of the solar polar regions, to advance our understanding of solar variability.
The highest science objective of SOLAR-C Plan A is measuring subsurface flow at high-latitude
regions, by local helioseismology techniques, based on Dopplergrams acquired by an HMI-MDI
type instrument. From this measurement we will derive differential rotation, meridional circulations
and convective flows in the upper convection zone. These flow-field measurements will then be
cross-correlated with surface magnetic field measurement, to reveal how magnetic flux istransported to the polar region, and how the polarity reversals take place as interplay between
plasma flows and magnetic fields. The predominantly vertical kG-field patches that Hinode has
found in the polar regions are large enough to be observed. A serious attempt will also be made forstereohelioseismology, for investigating deeper layers, including the best ever shot of the solar
tachocline region.
Solar irradiance measurements at various latitudes will, for the first time, enable us to measure theanisotropy of the total solar irradiance. The unexplained low photometric variability of the sun may
be explained by higher variability at higher latitudes, likely caused by faculae. If it is not the case,then we must conclude that the Sun is a rather atypical star, which will lead to more fundamental
questions in astrophysics.
The solar polar region also hosts various activity phenomena, and EUV and X-ray observations will
be important in studying them, in particular atmospheric plasma flows in solar-wind source regions.
Other scientific topics that SOLAR-C Plan A may address, exploiting its unique vantage point,
include evolution of heliospheric structures (by a heliospheric imager) and cosmic-ray transport in
the inner heliosphere (by a particle counter) when the optional payload is adopted..
Scientific Strategy
The outstanding issues confronting our current understanding of the solar dynamo may besummarized through several key scientific questions:
Q1) How is the global, cyclic, solar magnetic field generated?
Q2) What is the nature of flows in the polar regions of the Sun and how do they interact withmagnetism?
6Q3) How does the radiative energy output of the Sun depend on latitude?
Q4) How does magnetic activity shape the structure and evolution of the polar solar corona andhow does it affect the Earth?
Progress on these scientific questions requires detailed observations of the solar polar regions,where data is currently scarce and where much of the subtle interplay between plasma flows andmagnetic fields that gives rise to cyclic polarity reversals is thought to occur. The out-of-ecliptic
observations of the Sun, for the first time, will provide an opportunity for detailed investigations of
the magnetic structure and dynamics of the polar regions. High-latitude photospheric observations
will also provide an unprecedented vantage point for helioseismic imaging that can be used toprobe flows and fields in the deep convection zone and tachocline where solar activity is ultimately
thought to originate.
In addition to measurements at the photospheric level, the structures of the outer solar atmospherein polar regions and the heliospheric structures merit observations from outside the ecliptic. The
poles of sun undergo dramatic change during the 11-year solar cycle, driven by the dynamo actionin the solar convection zone. The polar vantage point gives unique opportunities for understanding
the origin of the fast solar wind spectroscopically and for stereo viewing of surface vector magnetic
fields, coronal structures, and Earth-directed CMEs in coordination with observatories near the
Earth. The unique inclination for the SOLAR-C Plan-A observatory will also permit unprecedented
measurements of the total solar irradiance. This may help resolve the discrepancy of the cycle
variation of the solar irradiance of ~0.1% while solar analogs vary, on average, by 0.3%.
With this in mind, we propose the following prime measurement targets for the SOLAR-C Plan-A
mission:
T1) Photospheric magnetic flux distribution and evolution in the polar regions
T2) Dynamical coupling between magnetic fields and flows
T3) High-precision measurement of total solar irradiance
T4) Structure and dynamics of the high-latitude solar corona and solar wind
The methodology, significance and expected scientific impact of each of these observational targets
are discussed in Section 2.1.
Measurement Requirements
The following observables are required to address the SOLAR-C Plan-A top science objectives:(A) full-Sun photospheric line-of-sight Dopplergrams for measuring the subsurface flows by
helioseismology, and (B) full-Sun photospheric magnetograms for tracking the evolution ofmagnetic fields on the surface.
These measurements must come from an orbit inclined to the solar ecliptic, with maximum orbit
inclination 40. The orbit must allow observations from >30inclination for periods of >40 daysfor each polar passage. To maximize the telemetry available, the orbit shall be circular with a 1year period, to synchronize the orbital motion of the spacecraft with that of the Earth.

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The orbit gives a unique vantage point for the other science objectives. The following observables
shall be taken to address them: (C) total solar irradiance, (D) full-Sun chromospheric images, (E)full-Sun transition region (TR) images, (F) full-Sun coronal images to monitor the dynamic activity
and the evolution of high-latitude structures, and (G) emission-line imaging spectra inchromospheric, TR, and coronal lines for investigating the source region of fast solar wind anddynamics of the outer solar atmosphere, (H) visible-light images monitoring interplanetary space
between the Sun and Earth, and (I) in-situ measurements including magnetic field, solar windprotons and electrons (TBD). Currently, (H) and (I) are treated as options in this Interim Report.
Photospheric magnetograms and Dopplergrams are made from multiple images. Not all the images
can be transferred to the ground due to the expected telemetry rate, so onboard data processing andcompression are mandatory to reduce the total data volume transferred, as was done on
SOHO/MDI. The field of view for imaging observations needs to cover the full Sun with sufficientspatial sampling for each observable. The science requirements for image size, cadence andduration for the helioseismic observables are discussed in detail in Section 2.1. Table 3 shows
estimates of the required data rates. A total average data rate of ~100 kbps is required.
SOLAR-C Plan-A Payload
The SOLAR-C Plan-A has been planned by attaching weight to science topics that are studied byremote-sensing instruments. The best condition for remote-sensing observations from the out-ofecliptic orbit is selected under the restriction of current or near-future technology.
The science payload to satisfy the primary measurement requirements consists of:- A visible light imager that can measure the full-Sun photospheric magnetic fields and line-ofsight Doppler velocity, similar to MDI on SOHO or a simplified version of HMI on SDO.
- A total irradiance monitor that measures the irradiance of the Sun. Multiple cavity monitors
are needed for self calibration.
- A light-weight EUV imager that monitors the transition-region and coronal activity in thepolar region.
- An EUV scanning spectrometer that measures the flow structures in the polar region.
The following instruments provide additional measurements of T4 and are listed as optional:- A heliospheric imager that observes the space between the Sun and Earth- In-situ instruments such as instruments measuring the solar wind and cosmic-ray particle
sampler (TBD)

Trajectory and Orbit
The major mission requirement of SOLAR-C Plan A is to observe the Sun from high latitude, withtarget specified as 40°. In order to achieve this target, the SOLAR-C Plan-A Working Group hasstudied a number of possible mission designs using various trajectory manipulation techniques. The
items considered are the tilt of the solar equatorial plane to the ecliptic, launcher capacity, planetarygravity assists, and the use of a highly efficient propulsion system. The launcher assumed is theJapanese H-IIA launch vehicle equipped with a solid motor upper stage. There are two possible
options: one is to use solar electric propulsion assisted by an Earth swing-by (SEP option), and the8other is to use both Jupiter and Earth swing-bys (Jupiter option). The latter is a purely ballistic
trajectory. The final orbit is nearly the same, a circular orbit with 1 AU distance from the Sun and 1
year orbital period, with a different time to reach the maximum inclination (5 years for SEP option
and 7 years for Jupiter option). If there is a strong requirement that the spacecraft has to be in the
final orbit before the solar maximum around 2024 (assuming launch in 2019), the SEP option is the
only currently available way to satisfy the time constraint. The orbit profile in the cruise phase isquite different between SEP and Jupiter options. In the first three years, no observations are
expected for the SEP option and there may be a short observing period of about a month at eachhemisphere in the following years until the spacecraft enters into the final orbit. On the other hand,observations are possible at any time outside the swing-by operation for Jupiter option, but hightelemetry is limited to positions near the Earth until reaching the final orbit.
Spacecraft
The SEP option requires a heavy propulsion system. The assumed μ20 ion engine system is an
upgraded version of μ10 that was used in JAXA HAYABUSA sample return mission. The newengine itself has undergone endurance testing in the lab but no flight heritage yet. The μ20 ion
engine system consumes 6 kW power, and a lightweight solar array paddle with large area isnecessary. The total wet weight of the spacecraft in the SEP option is ~1,200 kg with payload massof 130 kg. The total wet weight of spacecraft in the Jupiter option is ~750 kg with payload mass of
130 kg.
Launch vehicle
A single spacecraft launch with the JAXA H-II A rocket 202A is assumed. Some orbit cruise
methodologies require a kick motor and it becomes one of new development items in the Japanesespace agency after the interplanetary missions launched by M-III or M-V solid boosters. The use ofa kick motor is assumed for both SEP and Jupiter options. There is a solution in a SEP option by anorbit cruise via Venus swing-by without having a kick motor by reducing the launch opportunity.
The spacecraft attitude is three-axis stabilized to meet the requirements of the imaging instruments.
Angular momentum management occurs daily to weekly, using chemical thrusters. 300 kbps Xband downlink telemetry rate and 8 hr downlink time per day are assumed for an average data rateof 100 kbps, at a spacecraft distance at 0.56 AU from the Earth. Downlink stations are needed in
the northern and southern hemispheres on Earth. 1 Mbps Ka-band telemetry rate at the same
distance is under consideration to enable a greater telemetry volume.

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1. SOLAR-C Mission Science Goals
1.1 Origin of Solar Magnetism
1.1.1 Solar Magnetic Cycle
The solar magnetic cycle is one of the greatest puzzles in astrophysics. Similar magnetic cycles
have been observed in other stars; and the turbulent dynamo that controls these cycles is a key tounderstanding the cosmic magnetism. It has been known for centuries that the number of sunspotschanges quasi-periodically with 11-year cycle, forming the famous “butterfly” diagram (Figure 1
and Figure 2).

Figure 1: (left): Sunspot numbers and sunspot butterfly diagram (NAOJ)
Figure 2: (right): Averaged surface magnetic field strength (courtesy of D. H. Hathaway)

Many recent solar dynamo models are based on the so-called “flux-transport” paradigm whereby
the advection of magnetic flux by the axisymmetric flow in the meridional plane regulates theperiod and amplitude of the activity cycle. Most flux-transport models attribute the generation ofmean poloidal field to the emergence and subsequent dispersal of photospheric active regions by
the combination of poleward circulation and turbulent diffusion in the upper convection zone,
known as the “Babcock-Leighton” mechanism. The reversal of the global poloidal field occurs as
opposite-signed field from lower latitudes converges on the polar regions. However, how thisprocess couples to the dynamics and magnetic topology of the deeper convection zone is currently
unknown. Recent Hinode observations reveal the polar field is highly structured, with magnetic
field concentrations reaching 1kG, which is not captured by mean-field models. How the polar“magnetic landscape” changes with the solar cycle is one of the most critical questions forunderstanding the nature of solar magnetism. The magnetic field evolution is intimately coupled to
the flow field, including convection, differential rotation, and the mean meridional flow. In
particular, the latitudinal bands of photospheric magnetic activity reflected in the solar butterfly
diagram (Figure 1 and Figure 2) are closely linked to systematic variations in the differential
rotation known as “torsional oscillations” (Figure 3). A surprising result from helioseismology isthat the depth and strength of these flows increase in the near-polar regions. Furthermore, models
of surface flux transport indicate that the polar field strength is sensitive to variations in theamplitude and structure of the high-latitude meridional flow. However, reliable measurements ofthe magnetic and flow fields in the near-polar regions are currently lacking.
10There is a huge gap in our knowledge of the polar magnetism and dynamics, preventing us fromunderstanding the basic mechanisms of solar magnetic activity, and developing physics-based
forecasts of the solar activity and cycles.

Figure 3: Systematic variations in angular velocity known as torsional oscillations, illustrating theirlatitudinal and temporal variation several Mm below the photosphere (left) and their subsurface structurein the meridional plane (right; figure adapted from Vorontsov et al. 2002).

1.1.2 Generation and Emergence of Magnetic Field
In astrophysical objects, dynamo action can exist in plasmas with a seed magnetic field and flow
fields. However, sufficient conditions for dynamos are not well-determined. For solar and stellar
physics it is particularly important that dynamo processes can result in a cyclic behavior. Meanfield MHD theories of solar and stellar dynamos predict the cyclic behavior, which resembles theobserved properties such as the butterfly diagram for sunspot formation zone and polar fieldpolarity reversals. However, our understanding of the underlying physical processes is still
schematic Parker’s standard α-Ω mechanism that has been applied to the Sun and a wide range ofother astrophysical objects, stars, galaxies, interstellar medium.
The current global dynamo models involve several building blocks mainly known by theirrepresentation in mean-field electrodynamics: One is the α-effect that generates the poloidal fieldfrom the mean poloidal or toroidal field by a turbulent cyclonic flows (Item 1 in Figure 4), or by adisintegration of the surface active region fields. (Items 1’, 4, and 5 in Figure 4). The other is the Ωeffect that generates the toroidal field from the mean poloidal field by shear motions in the
differential rotation. These two effects couple with each other to re-generate the toroidal andpoloidal fields in repetition. The locations of each effect are now considered to be separated in theinterior, so that , in addition to these field stretching mechanisms, the transport of the fields are also
important ingredients: (Items 3 and 3’ in Figure 4). The three different transport mechanisms
generally considered are active transport toward the surface by means of magnetic buoyancy (risingflux tubes, formation of active regions), turbulent transport in the convection zone (turbulent
pumping, turbulent diffusion) and transport by the large scale meridional flow (flux transport
dynamo; Figure 5).

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0: Turbulent convection (plumes); 1: Generation/self-induction of poloidal B field (“-effect”) or 1’: Tilt of active
region; 2: Turbulent pumping of B field in tachocline or transport of B field by meridional flows from CZ into
tachocline (single for multi-cells flow?); 3: Field ordering into toroidal structures by large-scale radial and latitudinal
shear in tachocline (“-effect”); 3’: Surface shear layer or subsurface weather; 4: Toroidal field becomes unstable tom=1 or 2 longitudinal instability (Parker instability); 5:Rise and rotation of twisted toroidal structures; 6: Recycling of
weak field in CZ; 7: Emergence of bipolar structures at the surface; 8: Internal waves propagating in RZ and possibly
extracting angular momentum; 9: Interaction between dynamo induced field and inner field in the tachocline alongwith shear, turbulence, waves, etc.; 10: Instability of inner field and shearing via -effect at nuclear core edge? Is therea dynamo loop realized in RZ?

Figure 4: Solar internal dynamics and dynamo

12Figure 5: Flux transport dynamo (adapted from Dikpati and Gilman 2006 )The success of this beautiful theory hinges on the turbulent properties, magnetic and kinetic
helicities and diffusivity. However, direct numerical simulations, developed during the past decade,revealed significant limitations, such as the catastrophic helicity quenching, which severely restricts
the magnetic field growth. The potential solution is in studying the helicity balance including largescale circulations and helicity loss through coronal mass ejections. This study requires detailed
observations of the solar interior, magnetic field and coronal dynamics over the whole solar cycle.
It is particularly important to understand the mechanism of the cyclic polar field reversals revealed
in synoptic ground-based observations during the past three sunspot cycles. Because the Sun’s axis
is almost perpendicular to the ecliptic our knowledge of the polar magnetic field structure anddynamics is very poor.
The best opportunity to gain this knowledge can be provided by the out-of-ecliptic SOLAR-C
mission. It will provide an important insight into the structure of polar magnetic fields and themechanism of polarity reversals. The polar fields are largely unipolar, and it was believed that thepolarity reversals result from a diffusion or circulation process of magnetic flux transport from thelow-latitude zone where the flux emerges. However, recent high-resolution observations from
Hinode showed that the polar fields are highly structured, and, while the mean polar field is only a
few Gauss, the field in these elements is very strong. This discovery challenges the flux transportmodels, and opens a new opportunity of studying the relationships between the structuring andglobal field reversals. The out-of-ecliptic mission will also improve our knowledge of the largescale convection, differential rotation and meridional circulation, which are the key ingredients ofdynamo models.
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In terms of magnetic-field observations, compared to ultra high-resolution vector-magnetograms
obtained in the ecliptic plane, SOLAR-C has advantages in longer observation periods, more global
view of the polar regions, and capability of providing simultaneous flow-field observation in andbelow the photosphere, to facilitate study of interaction between plasma flow and magnetic fields.
Potential disadvantages are that part of the polar regions that lie on the Earth-side may be observedwith higher resolution from the ecliptic plane, and the lack of vector magnetogram capability fromthe spectro-polarimetry by an imaging polarimeter. If, however, as Hinode is finding out, the polar
fields primarily comprise relatively large, predominantly vertical kG-field patches, thesedisadvantages may not be so important.

1.1.3 Dynamics of the Solar Convection Zone and Tachocline
The differential rotation, meridional circulation, turbulent α-effect, and turbulent transport terms(turbulent diffusion and magnetic pumping) that form the basis of the solar dynamo models
ultimately arise from turbulent solar convection. Stellar observations confirm that convection
breeds magnetism; late-type stars such as the Sun with convective outer envelopes exhibit vibrantmagnetic activity whereas more massive stars with convectively stable envelopes are generally lessactive. Thus, understanding solar convection and the mean flows it generates is an essential
prerequisite to understanding solar magnetism.
Determination of the solar internal rotation profile and near-surface meridional circulation bymeans of global and local helioseismology, in the past three decades, has revolutionized solar
dynamo theory. Rotational shear has long been an essential ingredient in all solar dynamo models,as the principle generation mechanism for the global-scale toroidal magnetic flux (Ω-effect) that
ultimately emerges from the solar interior to form active regions. More recent flux-transport
dynamo models have further identified flux advection by the mean meridional circulation as a key
factor in the establishment of cyclic magnetic activity. In particular, the poleward advection of
emergent toroidal flux from lower latitudes may account for the polarity reversals of the polar field
and thus the cyclic reversals of the global dipole moment. Possible correlations between the highlatitude meridional flow speed and activity patterns such as cycle amplitude and duration arenecessary to distinguish among various dynamo paradigms. Regardless of the dynamo mechanism,
surface flux transport models indicate that the strength and distribution of the polar magnetic field(crucial for coupling to the heliosphere) is sensitive to the amplitude and structure of the highlatitude meridional circulation. Determination of the differential rotation and the meridional
circulation in the polar regions will thus provide unprecedented insights into the dynamics of theconvection zone, the operation of the solar dynamo, and the solar-terrestrial interaction, bringing
the helioseismology revolution to its ultimate fruition.
Observations of the solar photosphere reveal a hierarchy of convective motions, from solar
granulation (L~1 Mm) to supergranulation (L~35 Mm). Deeper in the convection zone convective
length and time scales increase as a consequence of the larger density and pressure scale heights.
Magnetic flux transport by supergranulation and granulation may contribute to the cyclic reversalof the large-scale poloidal field by working in concert with the meridional circulation. Not only
does convection influence magnetism, but the converse is also true; magnetism can influence thestructure of convection and investigating the nature of this nonlinear feedback provides valuable
insight into both phenomena. In particular, the structure and evolution of supergranulation in polar
coronal holes, where a net unipolar flux permeates the photosphere, may be significantly different
14than that at lower latitudes where the flux distribution exhibits mixed polarity. The Lorentz force
tends to decrease convective length scales but magnetically-induced enhancements in radiative
cooling may counteract this effect. Careful observations over extended time intervals (at least
hundreds of days for reliable statistics) at high latitudes are needed to elucidate the subtle nonlinear
feedbacks between solar convection, magnetism, and radiation.
Observational signatures of giant cells are notoriously difficult to glean from photosphericobservations but the unique high-latitude vantage point offered by SOLAR-C will provide newopportunities. The maintenance of mean flows by global convective motions is expected to produce
thermal gradients between the equator and pole that may be detectable by helioseismic inversions
or photospheric irradiance measurements (Section 1.2.3). Furthermore, theoretical and numericalmodels predict a change in morphology between global convective motions at high and low
latitudes in rotating spherical shells. The transition between polar and equatorial convective modes
occurs near the so-called tangent cylinder, a cylindrical surface aligned with the rotation axis andtangent to the base of the convection zone. This tangent cylinder intersects with the solar surface at
about 45 latitude. Possible changes in the subsurface flow fields inferred from local helioseismic
inversions inside and outside the tangent cylinder may provide a valuable and previously
unexploited observational signature of the elusive but extremely important giant cells. Hinodemeasurements have already revealed a systematic high-latitude alignment of supergranules that
may reflect the underlying influence of giant cells. Polar observations may also reveal otherphemonema such as inertial waves or precessing convective modes with longitudinal wavenumber
m=1, as seen in some numerical simulations.
Convective and meridional flows that converge toward the poles are accelerated in a prograde senseby means of the Coriolis force which can in some circumstances establish a cyclonic polar vortex.
The presence of a polar vortex remains controversial in current helioseismic inversions but isexhibited by some numerical simulations. Coordinated helioseismic determinations of thedifferential rotation and meridional flow at high latitudes will settle this issue, at least in the upper
convection zone, and will provide important constraints to global convection simulations,enhancing our understanding of angular momentum transport.
Like the polar regions, the base of the solar convection zone is an unexplored frontier of essential
importance to the operation of the solar dynamo. This is where photospheric active regions
originate, as the strong shear, convective pumping, and mild stratification promote the generationof coherent toroidal magnetic flux structures which buoyantly destabilize and rise. Flux-transport
dynamo models attribute the duration of the solar cycle to the equatorward flow speed near the baseof the convection zone. Furthermore, recent theoretical and numerical models indicate that thedifferential rotation profile throughout the solar convection zone is influenced by thermal andmechanical coupling to the tachocline. SOLAR-C will enable fundamental breakthroughs in our
understanding of solar internal dynamics and magnetism by probing the base of the convection
zone through the innovative and powerful technique of stereoscopic helioseismology (Section
1.1.4).

1.1.4 Probing the Solar Interior
In our effort to understand the mechanism behind the solar activity cycle, it is important that westudy the solar interior observationally. The only method known to us is helioseismology, which
15enables us to measure thermal structure and flows in the solar interior. In this subsection a very
brief introduction to helioseismology is given, to provide background for further discussions in
Section 2.
The solar 5-minute oscillations, discovered by Leighton et al (1962), were later identified as
manifestation of the global acoustic eigenoscillations, excited by acoustic emission by turbulent
convective motions. These eigenmodes are results of acoustic waves generated in the upper
convection zone travelling around the sun, some of which interfere with themselves constructively
and become resonant modes, whose properties are sensitive to the structure of the sun (Figure 6).
This motivated effort to develop methods of probing the solar interior based on precisely measured
eigenfrequencies.
In helioseismology, we start from observing surface wavefield, preferably by Doppler velocity
measurement, which is known to suffer less noise compared to intensity fluctuation measurement(by a factor of 5 or more is not uncommon). The wavefield is then characterized by Fourier (intemporal domain)/Spherical Harmonic (in spatial domain) analyses, to produce eigenfrequencies as
the primary product. The eigenfrequencies are then analyzed, typically by inversion methods, toestimate various quantities that affect wave propagation. One good example is adiabatic soundspeed distribution, and another is differential rotation in the sun (Figure 7), which is now measured
in low- to mid-latitude convection zone and deeper into the radiative zone, but not at high-latitudes
and in the central region of the sun. The limitations are mainly coming from an insufficient numberof eigenmodes that sample these regions. The profile of the solar differential rotation depicted in
Figure 7 gave observational constraints to solar dynamo theory. In fact, it contradicted columnar
rotation profiles which were generally believed to be the case for the sun, and understanding ofdynamo were based on.

Figure 6. (left): Acoustic ray paths for various modes. Each of these paths, after a trip around the sun, eventually
overlap with itself, and if it is found to be constructively interfering with itself, it becomes a resonant eigenmode.
Different modes have different paths, and therefore sample different parts of the sun.
Figure 7. (right): Differential rotation of the sun measured by (global) helioseismology (adapted from Thompson
et al 1996).

16

Figure 8. (left): Meridional flow speed measured by local helioseismology (adapted from Giles et al 1997)
Figure 9. (right): Subsurface local flows measured by a local helioseismology technique (adapted from
Haber et al 2002).

Building on these successes, in recent years, a new subfield of local helioseismology is developing.
As in the old version of helioseismology (often called global helioseismology), local
helioseismology does start from measuring wavefields. In local helioseismology, however, we donot measure (global) eigenfrequencies. Instead, we measure those quantities that characterize local
propagation of the waves, such as wave travel times between a pair of surface points (for example,
two bouncing points along a ray path in Figure 6), power distribution (in wavenumber-frequency
space) of local wavefield, etc. Unlike in global helioseismology, these measurements provide bases
for probing localized and asymmetric structure of the sun, not constrained by the symmetric mannerin which global modes sample the sun, as signified by symmetries of eigenfunctions.
Meridional circulation has thus been measured (Figure 8) by a time-distance method, and solar
subsurface ‘weather’ patterns (local flowfields) are now routinely measured by local power
spectrum analyses (Figure 9). Once again, the high-latitude regions are missing as, once beyond
about 60 off the disc centre, foreshortening and projection degrade the local wavefieldmeasurement significantly, although these parts are of great interest. Deeper layers are also difficult
to access, because of the following. As we see in Figure 6, probing a deeper layer requires
measuring wavefields at (at least) two regions on the sun that are apart by a great distance, in whichcase at least one, possibly both, of the region tends to be rather close to the limb, and therefore
suffers from projection and foreshortening. One possible way to resolve this difficulty isstereohelioseismology, in which multiple helioseismic observations from different view angles areused to observe wavefields without significant degradation due to projection and foreshortening.

1.1.5 Prediction of Solar Cycles
Prediction of the solar sunspot cycles is not only of great practical importance, it also represents anultimate test for our understanding of solar magnetism. So far, our predictions based on the α-Ωdynamo were not successful even with the input from helioseismology. The most successful
prediction of the previous solar cycle was made using measurements of polar magnetic field of theSun (Figure 10). The polar magnetic field reaches the maximum strength during the sunspot
minima, and it correlates very well with the maximum sunspot number of the following solar cycle.
In addition, it has been established long ago that a close correlation exists between the strength of17

geomagnetic indexes measured during the solar minima with the following sunspot maxima. This
correlation is probably caused by the high-speed solar wind from the polar regions.

Figure 10. Evolution of solar polar magnetic fields (solid: north and dashed: -south)

However, how the polar magnetic field strength is translated into the future sunspot number iscompletely unknown. It has been suggested that this may correspond to the transformation of the
global dipole magnetic field into the toroidal magnetic field of sunspots by the differential rotation(Ω-mechanism of the dynamo). However, the physics of this transformation and also themechanism, which determines the polar magnetic field strength, are unknown.
Thus, for developing physics-based forecast of solar activity it is extremely important to investigate
the structure and evolution of the polar magnetic fields, mechanisms of the magnetic flux transport
and interaction between the old and new magnetic fluxes. The SOLAR-C observations will provide
direct measurements of magnetic fields and flows, which will be incorporated into the dynamo
models and data assimilation procedures for predicting the solar magnetic cycles. In addition, it iswell-known that the solar 11-year cycles have a complicated structure with periods of very highactivity. It will be very important to investigate the correlation between these periods and thepreceding polar dynamics.

1.1.6 Scientific Questions to be Solved by SOLAR-C
We set the following four fundamental questions toward the understanding the solar magnetism and11-year activity cycle related issues that are tackled by SOLAR-C.
Q1) How is the global, cyclic, solar magnetic field generated?
Q2) What is the nature of flows in the polar regions of the Sun and how do they interact withmagnetism?
Q3) How does the radiative energy output of the Sun depend on latitude?
Q4) How does magnetic activity shape the structure and evolution of the polar solar corona andhow does it affect the Earth?
The relation between these questions and what targets should be observed with the science payload
is shown in Table 1. The detailed description on the background is shown in the following sections.
18

19

Table 1: Summary of scientific questions, targets and the required observations.

1.2 Exploration of Solar Interior and Solar Magnetic Activity
The primary goal of the SOLAR-C mission is to understand the origin of solar magnetic activity byobserving the magnetism and dynamics of the polar regions that are currently inaccessible for direct
observation. In Section 1.1 we identified four key scientific questions tackled by SOLAR-C. They
are summarized in Table 1. Here we describe each of the observational targets in turn and what
scientific impact they are likely to have.

1.2.1 Q1: How is the global, cyclic, solar magnetic field generated?
Understanding the origin, distribution, topology, and evolution of magnetic flux near the solar
poles is essential for understanding the origins of cyclic magnetic activity, for predicting theamplitude and duration of future cycles, and for assessing the impact of solar variability on the
Earth. Coupling between magnetic fields and flows in the polar regions governs the polarity
reversal of the global dipole moment and shapes the magnetic topology of the overlying corona.
The magnetic flux emanating from polar coronal holes fills the heliosphere and accelerates the fastcomponent of the solar wind, thus regulating the interaction between the Sun and planetary spaceenvironment. Further progress requires high-latitude magnetograms, Dopplergrams, andhelioseismic imaging.
Synoptic observations of solar magnetic fields have shown evidence that the magnetic flux at thefollowing (westward) edge of bipolar active regions is systematically transported toward the poles,
where it encounters poloidal flux of the opposite sign generated by the previous cycle. As lowlatitude flux continues to accumulate, the magnetic polarity of the polar regions reverses, leading toa corresponding reversal in the global dipole moment. However, how the process of the polarity
reversal actually occurs in detail and how it couples to the deeper convection zone is unknown.
What processes determine the distribution of surface and subsurface polar flux before and after areversal? This fundamental issue can and will be resolved by means of a magnetograph viewing thepole from an inclination much greater than the 7 provided by the ecliptic.
High-latitude magnetic field measurements are also needed to quantify the energy and flux budgetof the dynamo, the multipolar structure of the mean poloidal field, and the efficiency of poloidal
field generation by means of the Babcock-Leighton mechanism. Evidence for non-axisymmetricstructure in large-scale fields with low longitudinal wavenumber (m=1-3) might reflect thepresence of MHD shear instabilities. Measurements of the high-latitude magnetic field in thesolar photosphere are of fundamental importance for the understanding of the magnetic topology ofthe solar corona and its dynamical coupling to the solar interior, particularly with regard to smallscale flux emergence and coronal heating.
A complete understanding of how a polar field reversal occurs and what physical processes giverise to the solar activity cycle can only be obtained by considering magnetic field measurements
within the context of the global and local flows that generate them. This is addressed in question
Q2, Section 1.2.2.

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1.2.2 Q2: What is the nature of the flows in the polar regions of the Sun and how dothey interact with magnetism?
As alluded to in Sections 1.1.2 the dynamical origins of solar magnetic activity can only be fully
investigated by considering correlations between multiple observables. The most important
correlations are clearly those between the flow and the magnetic fields. Magnetic induction by
convection, differential rotation, and meridional circulation is responsible for the magnetic activity
patterns we see and the back-reaction of the Lorentz force on mean and fluctuating flows provides
an important diagnostic of subsurface dynamics. Investigating such interactions lies at the heart ofnearly all of the scientific objectives of the SOLAR-C mission.
Differential rotation and meridional flow in polar regions
Details of differential rotation and meridional flow in polar regions are currently unknown.
Differential rotation is inferred from global acoustic oscillations, which are not very sensitive to thestructure in high latitudes and in addition they do not allow for a discrimination between
contributions from the southern and northern hemisphere. Meridional flows are inferred from
Doppler measurements as well as local helioseismic techniques. Both approaches yield only
reliable data to latitudes of about 60 since resolution is limited due to foreshortening and, in thecase of helioseismic measurements, the signal to noise ratio is reduced due to the line-of-sightprojection of the p-mode Doppler signal.
Existing helioseismic measurements of differential rotation suggest sharp spatial and temporalvariations, including the possibility of a subsurface poleward jet (polar jet), but these measurements
can be confirmed only by observing the polar regions directly, from an out-of-ecliptic orbit and by
local helioseismology, which is free from the constraints of global helioseismology. The zerothorder structure of the meridional flow is approximately described by a single cell pattern withpoleward flow near the surface (and maybe in equatorward at the bottom). An additional flow cellwith opposite direction in polar regions has been inferred from some measurements, but is still
strongly debated. In addition to their dynamo implications differential rotation and meridional
flow have compelling diagnostic value in understanding the dynamics of the deep solar convection
zone. Particularly significant are possible correlations between zonal flows, meridional flows,magnetic flux, and thermal variations.
The most well established signature of dynamical coupling between the zonal and meridional flows,magnetic fields, and thermal gradients is that of the solar torsional oscillations. These arealternating bands of east-west zonal flow that evolve in close correspondence with the solar activity
cycle (Figure 3). Two components are evident from the data, including a low-latitude branch that
propagates equatorward in conjunction with bands of magnetic activity and a high-latitude branchthat propagates poleward on a comparable time scale. The high-latitude branch is stronger anddeeper than the low-latitude branch and may arise from distinct dynamical influences. Correlated
meridional flows and thermal signatures are known for the low-latitude branch but are currently
undetectable for the high-latitude branch. Variations of the meridional flow with the solar cycle
induced by magnetic or thermal forcing also have important consequences for determining thestrength and distribution of the polar field and the timing of polarity reversals. Coordinatedobservations of the meridional flow and the magnetic field in the polar regions will provide
quantitative estimates of flux transport and will thus elucidate the physical mechanisms that
underlie cyclic solar activity.
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The vantage point of SOLAR-C will provide an unprecedented opportunity to observe suchcorrelations and in particular their phase relationships, providing important constraints totheoretical and numerical models of the solar convection zone and dynamo by exploring a part of
the Sun that is virtually unknown.
Tachocline dynamics/stereoscopic measurements
In order to observationally constrain the subsurface location of dynamo activity it is necessary todetect either field-induced perturbations of the sound speed or related flow fields. Sunspots andrelated bipolar active regions in the solar photosphere are thought to originate from the buoyant
destabilization and subsequent emergence of strong, coherent toroidal magnetic flux structures thatare generated deep in the solar interior, most likely in the tachocline or the lower convection zone.
Some theoretical and numerical models predict that such concentrated toroidal flux structures may
support coincident zonal jets of fluid as large as 100 m/s that help stabilize them against magnetic
tension and buoyancy. Although direct detection of such deep-seated toroidal magnetic flux
structures may be beyond the sensitivity of helioseismic inversion techniques, the zonal jets they
induce may be detectable by stereoscopic helioseismology in which coordinated observations using
SOLAR-C and another instrument such as HMI or GONG would provide the very long baselines
needed to measure the relatively long-wavelength modes that sample the deep interior with local
helioseismology.
Currently one of the most pressing issues concerns the strength of the subsurface toroidal fields that
ultimately emerge to produce photospheric active regions; current estimates range from 10 to 100
kG. The predicted amplitudes of associated zonal jets increase with the strength of the field, and arepotentially only detectable toward the higher end of the estimated range (100 kG). Detections oreven upper limits of this order for zonal flows and toroidal field strengths in the deep convection
zone and tachocline could have a dramatic impact on dynamo models.
Stereoscopic helioseismology will also allow us to improve measurements of the meridional flowin the deep convection zone. Such measurements are challenging but extremely important, since themeridional flow speed at the base of the convection zone largely sets the 22-year period of the solar
activity cycle, according to current flux-transport dynamo models. Given the paucity of current data,
any limit imposed on structures, speed, etc. by observations would be of great importance. For
example, determining the depth at which the flow shifts from poleward to equatorward would
improve current estimates of the speed of the return flow at the base of the convection zone basedon mass conservation. The density contrast of 100 between the lower (0.7R) and upper (0.96R)
convection zone implies values of order several m/s, which would be extremely difficult to measure
directly, even with a stereoscopic method.
In addition, the tachocline might also play a central role in determining the structure of differential
rotation throughout the convection zone (deviation from Taylor-Proudman state, in which therotation rate remains constant along the direction parallel to the rotation axis). Here additional
constraints on the thermal structure, especially latitudinal temperature variations are crucial.
Current attempts to measure such variations using the GONG and MDI data will be greatly
improved by supplemental data from the unique perspective of SOLAR-C. The required accuracy issubstantial, since models predict only a few Kelvin variation, as compared with a backgroundtemperature of about 2 million K at the base of the convection zone. Yet, if magnetic effects can be
22

properly taken into account, local helioseismic inversions and irradiance measurements (Sec. 1.2.3)
may be able to provide some estimate of the pole-equator temperature difference in the solar
surface layers, if not in the deep convection zone.
Stereoscopic techniques will also benefit global helioseismic inversions, which are currently
impeded by incomplete sampling of the solar surface. Observations from a single vantage point
only sample one side of the Sun rather than the entire spherical surface on which sphericalharmonic functions would form an orthogonal set. Multiple vantage points would greatly improve
spatial coverage. In case of SOLAR-C, with its position in the orbit around the Sun synchronouswith Earth, for maximizing the telemetry rate, the increase in spatial coverage is not so great but it
still offers an increase particularly when its heliographic latitude is high.
Determining the differential rotation, the meridional flow, and convective patterns in the upper
convection zone at high latitudes will be achieved by applying well-established local
helioseismology techniques previously confined to lower latitudes. Probing deeper layers will bedone by a novel technique of stereoscopic helioseismology from multiple vantage points.

1.2.3 Q3: How does the radiative energy output of the Sun depend on latitude?
The main goal of the proposed SOLAR-C Luminosity and Irradiance Monitor (LIM) is to address
the basic question, ‘How precisely do we know the luminosity of the Sun?’ A more detailed
description of the science goals can be summarized by the following four fundamental questions
related to the radiative output of the Sun:
• Are the solar irradiance variations observed from Earth due to flux redistribution in space or due
to solar luminosity variations?
• Is the radiant output at the poles different from the values we observe in the ecliptic?
• What is the latitudinal temperature gradient?
• Why is the variability in irradiance of the Sun a factor of three smaller than that of sun-like stars?
The first three questions arise because up to now, the irradiance was observed exclusively from thein-ecliptic vantage point. The fourth question was first asked by Lockwood et al. (1992) who notedthat compared with Sun-like stars, the irradiance variations induced by the solar cycle appear to besmaller than in Sun-like stars. It has been proposed that this exceptional property of the Sun is due
to our equator-on viewing angle (Lockwood et al. 1997, Radick 1998) but it is not yet clear whether
this explanation can fully account for the observed difference (Knaack et al. 2001).
The observational answers to these fundamental questions require measurements of the solar
radiative output at higher heliospheric latitudes than what an ecliptic-bound mission can provide.
Observations from the SOLAR-C Plan-A platform will allow us to verify the view-point hypothesis.
SOLAR-C Plan-A is the only planned mission that intends to fly a radiometer out of the ecliptic.
On average, stars are seen from a latitude of 30º. A trend in the solar irradiance level as SOLAR-C
LIM approaches this inclination will allow us to confirm or reject the view-point hypothesis. Thismeasurement will also be possible check the reliability of TSI models at the Maunder minimum erathat were reproduced using various indices with basic data of cosmogenic isotope that record thepast solar magnetic activity. This is of vital importance in the future TSI prediction for the Earth’s
environment.
23

Reproducing the solar luminosity is the ultimate requirement to all stellar evolution models. In
other words, the total solar irradiance is used to calibrate models of stellar evolution. The energy
production in the core of a main-sequence star is balanced by the radiant emission at the stellar
surface. Since the energy production rate as well as the radiant emission strongly depend on thetemperature of the core and surface, respectively, accurate measurements of the solar luminosity
provides a direct link to the temperature not only at the solar surface but also in the solar core.
Latitudinal irradiance scans also allow assessment of latitudinal temperature gradients which might
be linked to meridional flows.

1.2.4 Q4: How does magnetic activity shape the structure and evolution of the polar
solar corona and how does it affect the Earth?
In addition to measurements of photospheric and subsurface motions and magnetic fields, it is thefirst opportunity to observe the outer solar atmosphere of the polar region in detail.
Coronal Structures and Activity in Polar Regions
The Japanese solar physics community monitored the brightness of the solar corona at a height
above the solar limb by small ground-based coronagraphs for 60 years covering five 11-year solar
activity cycles (Figure 11). In the space program, the whole Sun was monitored in soft X-ray
observations by the 2nd Japanese solar mission, Yohkoh, from 1991 to 2001 (Figure 12).

Figure 11. Coronal brightness variation in time-latitude diagram from the observation in the coronal green
line (Fe XIV 5303) by the ground-based coronagraph at Norikura Solar Observatory (NAOJ)

24

Figure 12. Cyclic variation of coronal structures in soft X-rays from Yohkoh (ISAS/JAXA)

Similar to the butterfly diagram of sunspot and surface magnetic fields, long-term coronal activity
also shows the butterfly pattern in the time-latitude diagram. It has been well known since 1950’sthat there is a pole-ward migrating blanch in the diagram (Waldmeier 1957).
The polar regions of the Sun outside the solar maximum have a ray-like structure called ‘polar
plume.’ The structure is well observed by white-light observations at the eclipse time orcoronagraph observations in space. It is known from Skylab era that its electron temperature isslightly lower than the inter-plume region. The STEREO observations of polar regions with twoview angles show where polar plumes observed with two coronal EUV imagers are rooted.
However, the magnetic activity at the root could not be measured by a foreshortening effect byobserving the polar region from the ecliptic plane, so that how these structures are formed has notbeen fully understood. Imaging observations of polar plumes from a vantage point in the SOLAR-C
orbit will clearly address where and how they are formed. In addition to measurements ofphotospheric magnetic fields and imaging of transition region (TR) and corona, EUV imaging
spectroscopy is of vital importance in revealing what basic processes are occurring at the roots ofpolar plumes and how plasmas in plumes move as results of the processes. One of the importantmeasurements is to detect clear evidence of Alfvén waves in the polar plume or in inter-plumeregions as we stress in the next sub-section of solar winds.
Near the solar maximum phase, the polar region is covered by many closed magnetic fieldstructures with a ray-like shape at their tops, streamers. The phase of their appearance is near thetime of polar magnetic field reversals. The coronal imaging observations from the SOLAR-C outof-ecliptic orbit will clearly tell us how magnetic fields of these structures are connected at thephotosphere and how these structures are formed near the solar maximum phase.
In addition to characteristic global magnetic structures in polar regions, small bipoles are also
universally observed there. They were called X-ray bright points or coronal bright points that werefirst observed in an early sounding-rocket experiment with soft X-ray imaging telescopes. They are
distributed all over the latitudes. From photospheric and coronal observations near the disk center,
opposite magnetic polarities appear to be cancelling at the photospheric level beneath coronalbright points. It is believed that they are produced though magnetic reconnection in the corona by25

interaction between opposite magnetic polarities that have already seen above the photosphere, not
as a direct emergence of a small bipolar structure. The elemental process occurring up to the midlatitudes is being investigated by Hinode with its high-spatial resolution in a narrow field of view.
The small bipoles are replenished rapidly with time and the total magnetic flux exceeds thosebrought from the global magnetic activity showing a quasi-periodic activity of roughly 11 years.
The occurrence frequency of the bright points measured in low-latitude regions is found to beroughly constant in contrast to the normal solar magnetic activity that has a quasi-periodic 11-year
cycle. The low-amplitude or almost constant cyclic behavior may be the presence of a different
dynamo action that locally functions. It is not clear at present that the emergence rate of such
bipoles in polar regions is the same as that in low latitudes. A direct observation with lessforeshortening effect from the out-of-ecliptic orbit with less overlap by foreground and backgroundcorona will clearly answer the behavior of small bipoles in polar regions.
Hinode high-cadence X-ray observations have shown that polar regions are more dynamic than
expected from the extension of Yohkoh X-ray observations as revealed from a higher occurrence
frequency of coronal jets (Cirtain et al. 2007). Hinode has also revealed from a high-resolution
vector magnetic field measurement that the formation of coronal jets in a polar coronal hole isfound to be associated with newly discovered photospheric kG concentration of polar magnetic
fields through magnetic reconnection (Kamio et al. 2009). Such events may produce Alfvén waves
that are thought to be important to the acceleration of the fast solar wind. EUV spectroscopy withenhanced sensitivity will reveal the coronal dynamics in polar region.
Solar Wind
Supersonic plasmas flow into interplanetary space from the Sun as the solar wind. Many in-situ
observations in the ecliptic plane have shown that the speed of the solar wind ranges from 300 to
800 km/s. Timing analysis of the fast solar wind (600-800 km/s) with full-disk X-ray images andglobal open magnetic field structures of the Sun have revealed that polar coronal holes and thoseextended from polar regions to lower latitudes are the source of fast solar wind. The solar windspeed measured by Ulysses whose orbit is largely inclined to the ecliptic orbit clearly showed that
there is a latitudinal variation in the solar wind speed and that the speed is bi-modal in the solar
minimum phase, that is, the slow wind with outflow speed of 300-400 km/s in low latitudes and thefast wind of 600-800 km/s in high latitudes (McComas et al. 2003, 2008; Figure 13). It is also
known from Ulysses that the slow solar wind also appears in the high-latitude region near the solar
maximum phase (Figure 13). It is widely accepted that the fast solar wind is so fast that the speed
cannot be explained by Parker’s thermally driven wind model alone. An additional energy input
needs to be given to the plasma flow. The interplanetary scintillation (IPS) observations show thatthe acceleration of the solar wind is almost finished within 30 solar radii.

26

Figure 13. Solar wind velocity measurements by Ulysses (McComas et al. 2008)

There were two categorized candidates for the mechanisms of fast solar wind. One is a fast jet
produced at the transition region or low corona, and the other is the acceleration by waves in thecorona. The possibility of the former disappears after SOHO observations because the high-speed
jet is not universally observed in the EUV imaging and spectroscopic observations in the transition
region and low corona and because the acceleration of the flow speed at the location between the
low corona and a distance of 30 solar radii becomes mandatory from the tracking of white-lightcoronal structure seen in the coronagraph (Sheeley et al. 1997) and from the evaluation of windspeed by Doppler dimming technique (Cranmer et al. 1999). The unexpectedly larger number ofjets in polar coronal holes has been found by Hinode X-ray imaging observations (Cirtain et al.
2007; Figure 14). Although it may contribute to some fraction of the total mass with a million
degree for the fast solar wind, it cannot explain the acceleration of the solar wind in the most innerheliosphere within 30 solar radii.

Figure 14. North polar coronal hole observed with Hinode XRT showing a coronal X-ray jet

One of the most promising mechanisms is acceleration of the wind plasma by Alfvén waves. The
wave itself has normally been found in solar wind plasmas by in-situ measurements since the early
space age (Belcher and Davis 1971). The source of the Alfvén waves is not well understood. It may
be a wave simply produced by the photospheric motion of magnetic field structure by granulemotions or a wave created by magnetic reconnection in the low corona. A high-speed solar wind isreproduced in Suzuki & Inutsuka (2005) by the photospheric lateral motion of the magnetic field.
However, clear evidence of Alfvén wave with sufficient energy flux has not been found in the polar27

coronal hole. The wave signatures near the base of polar coronal holes that have been reported aremostly identified as slow-mode waves. One of the signatures for the presence of Alfvén wave near
the base of the corona is enhanced line broadening of the coronal emission lines that are always
observed in polar coronal holes. If it is found that this broadening is actually caused by Alfvén
wave, where the energy flux of Alfvén waves created and what amount of energy flux Alfvén
waves have could be understood. One way is to look at the polar region with an imaging
spectrometer from various angles or to look at the same point stereoscopically with two imaging
spectrometers that locates at different heliocentric latitudes. For this type of observation, higher
spatial and spectral resolutions than those realized as SOHO SUMER and Hinode EIS is not
required.
The solar wind is largely accelerated to the coronal sound speed near 2-3 solar radii and becomesupersonic beyond that point. In the upper part of the corona at a radial distance above ~2 solar
radii, the hydrogen Lyman alpha line (H I) at 1216 Å shows a spectrum of very wide line width.
This was first discovered by the sounding rocket experiment in l979 (Kohl et al. 1980). The
emission mechanism is the resonance scattering of photons nearby the Lyman alpha wavelength bya small amount of neutral hydrogen in the corona, and was identified in Gabriel (1971). A similarline profile is seen in ionized oxygen resonance lines O VI at 1032 Å and 1037 Å (Figure 15). The
number of scattered photons depends on the outflow speed of each hydrogen or ionized oxygen
atom that absorbs a photon and re-emits. Faster plasmas scatter the smaller number of photons(Doppler dimming). The bulk flow speed of the solar wind plasma was evaluated from the scatted
photons by SOHO UVCS in the inner corona that cannot be accessible by IPS observations. The
cause of the largely enhanced line broadening is not fully understood. One definite contribution isfrom the line-of-sight component of the solar wind bulk outflow and the other is from the local
mass motion of atoms. The local mass motion consists of the thermal Doppler motion and some
additional mechanical motions. The latter contribution contains information of the heating andacceleration of solar wind plasmas.

Figure 15. Enhanced emission-line broadening at observations above coronal holes with SOHO UVCS
(from Kohl et al. 2006).

Largely enhanced emission-line widths have been found in O VI observations with SPARTAN and
SOHO UVSC (Figure 15). The line width implies that there will be large ion mass motions28

perpendicular to the line-of-sight direction. It is sometime referred to be 100 MK at height above 2
solar radii, while the electron temperature at that height is about a million degree. The enhanced
kinetic temperature is interpreted as a perpendicular motion of ions to the magnetic field line (T),and its highly anisotropic kinetic temperature (T/T// ~100 >>1) is due to the heating by somemechanisms. As a potential mechanism, the ion cyclotron resonance heating has been proposed(McKenzie et al. 1995; Cranmer et al. 200?). The mechanism requires a high-frequency Alfvén
waves of 104 Hz in the corona, which cannot be expected by the low frequency (0.001-0.01 Hz)
photospheric granule motions. The estimate of T on O VI observations depends on the model
outflowing atmosphere in polar coronal holes. A solution to satisfy the UVCS observation ispossible for simple isotropic kinetic temperature (T=T//) cases (Raouafi and Solanki 2006) and for
reduced anisotropic (T// < T < 10 T//) cases (Nakagawa 2007). Observations of polar coronal holes
from different helio-latitude by the SOLAR-C Plan-A spacecraft from out-of-ecliptic orbit orstereoscopic measurements with a spacecraft in a low Earth orbit or at L1 Lagrangian point will
resolve this discrepancy and leads to correct understanding of motions of ionized atoms withdifferent mass from hydrogen. The global magnetic-field structure inferred from the SOLAR-C
magnetic field measurement and coronal imaging from a high-inclination vantage point will also
become a help for understanding the magnetic structures in polar coronal holes as a model input.
The solar wind speed in the polar region changes in the solar cycle as shown in Figure 13. When
the SOLAR-C spacecraft is in the final orbit, the spacecraft can see each polar region for a few
months per year in a good condition above 30 from the solar equatorial plane. The change of flow
speeds at the coronal base from the open magnetic-field region in each polar region can first be
monitored during the rapid change in the polar reversal phase. The flow speed as a function ofevaluated expansion factor of magnetic field structures in the polar region, which is considered to
be a key factor for the fast wind speed, may be observed by SOLAR-C.
Coronal Mass Ejections & DisturbanceCoronal Mass Ejections (CMEs) are the most powerful source of the heliospheric disturbances andthe primary cause of the geomagnetic storms. Understanding and predicting their initiation,
evolution and interaction with the planetary atmosphere is one of the most important topics in solar
physics and the space weather studies. The other source of geomagnetic disturbance is the Corotation Interaction Region (CIR), the large-scale plasma structures generated by the interaction of
fast and slow solar winds. The recurrent MeV-ion events and the decrease of the galactic andanomalous cosmic ray intensity are also associated with CIR.
Initiation of CMEs and acceleration of the solar wind in the lower corona are being investigated indetail by Hinode, using the combination of photospheric and chromospheric observations by the
SOT and coronal observations by the EIS and XRT. In order to investigate the later evolution of
CMEs and CIRs in the inner heliosphere, it is necessary to have continuous imaging observationthat covers from the outer corona to the interplanetary space.

29

Figure 16. (Left and Middle) Evolution of a CME event observed with STEREO/HI-1 (from Harrison et al.
2009). (Right) Overhead view of the CME travel in the Sun-Earth interplanetary space (from Reames
1999)

The STEREO Heliospheric Imager instruments clearly demonstrated that CMEs and CIRs near
1AU and beyond can be detected by such instruments (left and middle panels in Figure 16;
Harrison et al. 2009), and provided the first stereoscopic view of the solar corona and the innerheliosphere. With a STEREO/HI-like imaging instrument on board SOLAR-C/Plan-A, we will beable to obtain the first overhead view of the heliosphere, which has been often drawn as "cartoon
models" by many researchers (right panel in Figure 16). It will allow us, for the first time, toinvestigate the evolution of CMEs and CIRs in the background Parker-spiral solar wind.
Combination of the context heliospheric image from SOLAR-C/IHI and the in-situ plasmameasurement in the ecliptic by a spacecraft such as ACE, WIND, GEOTAIL and Cluster will be ofgreat synergy. For example, the basic structure of CME is believed to be twisted flux rope that iseither already formed before eruption or generated during the eruption by the post-eruptive
reconnection (Forbes 2000). However, it has been pointed out that the interplanetary CMEs
(ICME) observed by in-situ instruments at 1AU often show more complicated structure without
well-defined flux rope (Richardson and Cane 2004). With the overhead view of the CMEs fromSOLAR-C, we can tell where in the CME the in-situ instruments are observing, thus allowing theinterpretation of the in-situ measurements in the global context.
The overhead view also has a great impact from the space weather point of view, as we will be able
to see what part of the CMEs and CIRs are going to hit the Earth that allows more reliableprediction of the solar wind properties at the Earth.
Solar energetic particles (SEPs) associated with flares and CMEs and cause significant damage ofthe spacecrafts and disruptions of the human activities in the space. Moreover, acceleration ofenergetic particles is one of the most fundamental processes in the space and astrophysical plasmas.
Understanding the production of SEPs and their propagation to the Earth is therefore one of themost important subject in the solar and heliospheric physics from both practical and academic
points of view. Indeed, the solar corona, interplanetary space and the Earth's magnetosphereprovide powerful and complementary laboratories for particle acceleration. While imaging of thesolar corona shows the global change of the magnetic topology, in the magnetosphere, detailmicroscopic plasma processes can be observed by the in-situ measurements. On the other hand,with the out-of-ecliptic heliospheric imager and the other in-situ instruments, the interplanetary
space provides a unique laboratory in which one can see both the global dynamics and morphology
and the microscopic plasma processes.
30

Phenomenologically, SEP production can be classified into two categories, one associated withsolar flare in the low corona (or more physically, magnetic reconnection), and the other associatedwith CME-driven interplanetary shocks. SEPs are usually observed at 1AU after they escape fromthe acceleration site and propagate, primarily along the interplanetary magnetic field, to thespacecrafts. Note that shock acceleration also is believed to be occurring in many astrophysicalphenomena, such as generation of the galactic cosmic rays (GCRs) in the supernova shocks andgamma-ray bursts.
SOLAR-C/HI can address this issue by providing the global context of the interplanetary space inwhich acceleration and transportation occur. From the overhead view of the CMEs we can measure
the mach number and, with an assumption of Parker-spiral-like ambient solar wind, the angle
between the shock normal and the ambient magnetic field in the upstream. Also by assuming theParker spiral, one can determine the magnetic connection of the shock front and the spacecrafts
observing the SEPs in the ecliptic. Using this information and a model of SEP transportation, one
can determine where in the CME shock the SEPs come from, thus allowing the comparison of theacceleration models and observations.
In-situ Measurements in the Inner Heliosphere
The past solar activity has been investigated by cosmogenic isotopes such as 14C and 10Be. These
are created by the interaction between atoms in the Earth’s atmosphere and cosmic rays in a certain
energy range (~ a few GeV or below). These cosmic rays mostly arrive at Earth through the
heliosphere during a course of a long travel in our galaxy. Due to the change of magnetic fieldstructure in the heliosphere according to solar cycle variations, the number of the cosmic rays
suffers from significant modulation. The transport process of cosmic rays in the heliosphere was
first shed light on by Parker (1965). The theory describes the time evolution of the cosmic ray
distribution function by processes of (1) convection and adiabatic energy change by solar windvelocity, (2) turbulent irregularities, and (3) gradient, curvature, and current-sheet drift. The basic
formulation is still valid. However, there are uncertainties in various parameters.
Ulysses has executed the fist measurements of cosmic rays at high heliocentric latitude of 80. It
was expected that the cosmic ray intensity would increase an order of magnitude compared by themeasurement near the ecliptic plane because the local interstellar cosmic rays can travel to the innerheliosphere along the magnetic field without significant modulation. In contrast to the theoretical
expectation, Ulysses found that the cosmic rays increase only a few ten percents from ecliptic to thehigh latitude during the solar minimum. This indicates enhanced diffusion by turbulent
irregularities in the high latitude. Thus, the transport model of cosmic rays in the heliosphere isnot consistent with the Ulysses observations. The maximum inclination angle of 33 from theecliptic plane in the SOLAR-C Plan-A orbit is smaller than that of 80 in the Ulysses orbit.
However, a fast scan with a constant period of a half year at 1 AU distance from the Sun is anadvantage in the SOLAR-C Plan-A orbit for the scan range from the northern maximum tosouthern maximum latitude, and may have an advantage of investigating how the particle counts
vary with the solar activity cycle in a fixed observing condition. The measurement may produce
unique data set to enhance the value of long-term data on the variation of cosmogenic isotopes thatrecord the past solar activity. Without this measurement no progress is expected from observations
for the transport of cosmic rays in the heliosphere in the coming decades.
31

2. Scientific Requirements
2.1 Scientific Requirements
2.1.1 Requirements for Helioseismology Observables
The success of global helioseismology, which is centered at precise measurement of the solar
eigenfrequencies and inversion of these frequencies, is well documented. In particular, mapping thesolar differential rotation (Figure 7) and meridional flow (Figure 8) throughout the low- to midlatitude convection zone had a great impact on our understanding of the solar dynamo.
One inherent limitation of global helioseismology comes from inverse analysis based on the linear
perturbation theory, and the symmetry of the global eigenfunctions; the structure inversion is
primarily about estimating the spherically symmetric component of the thermal structure of the sun(even inversions for aspherical components of the thermal structure of the sun is accessing only
those components that are axisymmetric and north-south symmetric), and the rotation inversion isabout measuring the axisymmetric, north-south symmetric component of the solar internal flow.
Moreover, the sensitivity kernels for rotation inversion, derived from the eigenfunctions, do nothave large amplitude in higher latitude region.
Local helioseismology, based on local measurement of wave propagation, particularly wave travel
times, is better suited for revealing local structure free of any kind of symmetry that are mentioned
in the above. Local helioseismology has already been successful in measuring near-surface
component of meridional flow (Figure 8), which is the key ingredient in flux-transport dynamo, up
to around 60 latitude (Giles et al. 1997). Extending the measurement to higher latitudes is one of
the main scientific targets of SOLAR-C Plan-A.
Thus there are two main objectives of heliseismology observations from an inclined orbit: 1. Highlatitude local helioseismology: measure local wavefield in high-latitude region better than fromwithin the ecliptic place, for local-helioseismology analyses of various flows in the near-surface(and if possible, deeper) layers, such as meridional flow. 2. Multi vantage-point helioseismology: in
combination with observing facilities in the ecliptic plane, observe waves with large skip angles
better, for helioseismology analyses of deep interior, such as mid-latitude tachocline regions. The
first objective is more promising in producing important results, whereas the second objective ismore challenging.
Helioseismology requires long continuous observing periods at a high time cadence in order toresolve the various spatial and temporal scales associated with processes of generation anddissipation of solar magnetic fields. These scale range from supergranulation, which defines themagnetic network and is a primary source for magnetic field diffusion and has a typical spatial
scale of 30 Mm and lifetime of about 1 day, and larger-scale dynamics of active regions occupying
50-100 Mm and evolving on the scale of a week, to the global meridional and zonal flows
occupying the whole convection zone 200 Mm deep and varying with the solar cycle on the timescale of a year.
These three fundamental scales of solar dynamics and activity define three basic observing patterns
that are needed to achieve the primary objective above: relatively short (8 hour long) but highresolution (512512 pixels) time series to investigate the structure and dynamics on the32supergranular and magnetic network scales; longer (7-14 days) and high-resolution (512512
pixels) series to follow the evolution of individual active regions (particularly, on the far side of the
Sun in conjunction with Earth-side observations), and long time period (36-72 day) low resolution(256256 pixels) observations to investigate the global dynamics and circulation of the convection
zone down to the tachocline by both local and global helioseismology techniques.
All these observing programs require a 1 minute cadence for the Dopplergrams. However, a strict
temporal continuity requirement (95% coverage) is required for 8-hour chunks of these programs.
The longer observing programs may have gaps. However, the total coverage should be at least 7080%, in particular, for the global helioseismology 36-72-day runs in order to prevent the aliasingproblems in the measurements of mode frequencies and frequency splitting. The precision of thefrequency measurements is inverse proportional to the square root of the observing time. Currently,
72-day runs are used in the SOHO/MDI Structure program that is similar to the low-resolution(256256 pixels) SOLAR-C program. These runs provide inversion results for solar structure,asphericity and rotation through the whole convection zone, tachocline and outer radiative zone.
Shorter, 36-day runs will still provide new results for the convection zone sufficient for studying
the torsional oscillations and other global variations. However, the ability to resolve the tachocline
structure and dynamics will be greatly reduced. The investigation of the solar-cycle variation in thetachocline region may require 12-month long averaging of the frequency measurements. The
minimum science requirement for exploring the solar-cycle variations in the polar regions is tohave two or three 36-day runs per year.
Table 2. Helioseismology Observing Requirements
Measurement Objective(s)Length of

Image size

Imageobservation(pixels)

cadence

Global

36-72 days

256256

1 min

Local

8 hrs

512512

1 min

Meridional circulation, latitudinal andlongitudinal structures, secondary
cells,
relationshiptoactivelongitudes, magnetic flux transport

Local

8 hrs

512512

1 min

Supergranulation and large-scaleconvection patterns in polar regions(super-rotation, wave-like behavior,
network, flux transport, relationshipto coronal holes)

Local

8 hrs

512512

1 min

Structure and dynamics of the highlatitude tachocline (oblateness, flows)

Local

36 days256256

1 min

Tomography of the deep interior:

Local
36 days256256

1 min

Differential rotation and torsional
oscillations at high latitudes, subsurface polar jets

stereoscopic observations (SOLARC-SDO, SOLAR-C-GONG)

Technique(time-distance,
holography)

33

The different measurement objectives, the observing technique and the observing requirements arelisted in Table 2. Global helioseismology, which assumes symmetry about the rotation axis of theSun, requires a very accurate determination of the frequencies of the Sun’s acoustic modes of todetermine the longitudinal rotation as a function of depth; the longest observation times (36-72
days of continuous, 70-80% coverage data) are needed to determine the frequency and frequency
splitting of solar oscillation p-modes of low angular order, m. These modes, in particular, zonal
m=1 modes, propagate into the polar regions and provide the information about the polar structureand dynamics. The current ecliptic-viewpoint observations do not provide measurements
sufficiently accurate for helioseismic inversion of the solar structure and rotation in the polar
regions. It is important to complement the SOLAR-C observations with ecliptic-viewpoint data inorder to infer the structure and dynamics of the whole Sun, from the equator to the poles.
It can be seen from the table that some of these measurements require high latitude viewing formany days at a time. This was the driving requirement for choosing an orbit inclination with thehighest possible inclination. The high-latitude viewing will provide a unique opportunity to study
the dynamics of meridional flows and rotation in the polar regions and search for deep longitudinal
structures in the tachocline by local helioseismology. In addition, the high-latitude orbit will allow
us to obtain better coverage of the deep polar regions using the observing scheme with two vantage
points, SOLAR-C and ecliptic viewpoint helioseismology instrument.
The required telemetry for the helioseismic observations is case dependent. However, it is obviousthat the reduction of the data by on-board processing is mandatory. For a measurement of Doppler
velocity map with a simultaneous measurement of line-of-sight magnetic fields, two polarization
states of I +V and I-V need to be observed at four wavelength positions for a good estimate of theline-of-sight magnetic fields. This methodology was used in Hinode filtergram observations. The
eight images obtained at eight observing times can be reduced to two index maps by simple
arithmetic operations, each can be converted to Doppler velocity and line-of-sight magnetic field bya known non-linear function. If we need Doppler velocity maps of 512×512 imaging points every 1
min, ~13 kbit/s data rate is required for the data compression efficiency of 3 bit/s (512 ×512 pixel ×
3 bit/pixel /60 s =13,107 bit/s) that is achieved in Hinode. For the case of local helioseismology that
requires both a large field of view and high-spatial resolution, 1024×1024 imaging points may benecessary, and the required data rate becomes ~51 kbps ( 1024×1024 pixel × 3 bit/pixel /60 s=52,429 bit/s).

2.1.2 Requirements for Magnetic Fields Observables
Along with the Dopplergrams, full-disk line-of-sight magnetograms of 1024x1024 pixels will betaken every 5 (TBD) minutes to observe the response of the magnetic field to the surface, subsurface and interior flows. Together, these observations will be used to study phenomena such as
the evolution of active regions, flux transport and the solar cycle field reversal. SOLAR-C,
complemented by near-Earth magnetograph observations, will enable us to follow the evolution ofactive regions continuously for much longer than the half-solar rotation now possible. Full-sun
vector magnetic field observations shall be done every 60 min (TBD) or high-cadence observationsof vector magnetic fields shall be possible for a 4×4 (TBD) arcmin field of view.
The required data rate for line-of-sight magnetograms of 1024x1024 pixels of 5 minutes cadence isabout 10 kbps (1024 ×1024 pixel × 3 bit/pixel /300 s =10,486 bit/s). When four Stokes polarization34

states of I, Q, U, V are obtained at four wavelength positions, 16 imaging data need to be
downloaded to the ground. 2048×2048 pixel vector magnetograms every 2 hours requires ~27 kbps
(2048 ×2048 pixel × 3 bit/pixel × 16 images/7200 s =27,962 bit/s). For the active region monitor
by vector magnetic fields of 256×256 pixel imaging points, the data rate of 51 kbps (256×256 pixel× 3 bit/pixel × 16 images/60 s =52,429 bit/s) is required for cadence of 1 minutes.

2.1.3 Requirements for TSI Observations
The fastest approach to reach the science goals for TSI measurements is to compare a latitudinalscan of the solar irradiance to contemporaneous measurements from in-ecliptic spacecraft from thelong-term TSI monitoring programs, such as SORCE/TIM and PICARD/PREMOS or theirsuccessors. A slightly more demanding approach consists in a self-contained statistical analysis ofseveral latitudinal scans to detect latitudinal trends in irradiance. Both approaches would benefit
form observations at higher heliographic latitudes but extrapolation of trends in the SOHO/VIRGO
measurements, which cover ±7°, suggests that 30° is sufficient to answer the above questions.
The SOLAR-C TSI instrument produces the on-source and off-source irradiance data by theperiodic operation of a shutter located in front of one of cavities. The average and peak data
producing rates are ~2 kbits/s and ~4 kbits/s after the data compression. The full FOV is about 5
and there is a 15 arcmin pointing tolerance. This instrument is assumed to be continuously operated
in the SOLAR-C out-of-ecliptic orbit after the nominal operation. There may be a request of a
special operation to do an intentional offset pointing for instrument calibration.

2.1.4 Requirements for EUV imaging and Spectroscopy Observables
For the imaging of transition-region (TR) dynamics, a band at 304 Å is suitable. Polar plumes arevisible in a band at 171 Å and 195 Å. Hotter coronal plasma components in coronal bright points
and active regions are observed in a band at 335 Å. These are observed by a telescope consisting ofmultilayer mirrors that has a spatial resolution of ~2 arcsec by ~1 arcsec spatial sampling. The fieldof view covers the whole Sun when the spacecraft points to the disk center. One of these bands isobserved at a time. Wide slit observations of EUV scanning spectrometer may supplement theimaging of the TR and corona, but the images obtained with the wide slit are not sharp due to theblurring in the spectral direction. The wide slit observation cannot be used for the stereoscopic
observations with a similar telescope in other line-of-sight directions. In the case of selecting a
grazing-incidence telescope as a SOLAR-C coronal imager, observations of polar plumes may be
performed in a low-contrast condition and the TR observations become impossible. A nominal
operation is a combination of the full-disk observation at ~1 hr interval low cadence and smaller
field-of-view observations with a higher cadence of few min. A fast observation consuming all
allocated telemetry in a short time needs to be supported. ~20 kbps data recording rate is necessary
when a data compression of 1.5 bits/s is considered.
Emission line spectroscopy with focusing optics is required to measure flows of solar-wind source
region and to detect a signature of Alfvén waves that is a promising candidate to accelerate fastsolar wind. A candidate of wavelength is in EUV. Although EUV instruments are relatively heavy,
the SOLAR-C EUV scanning spectrometer needs to be achieved in a lighter weight condition35

because the total payload mass is limited. It is possible to enhance the instrument sensitivity of asmall EUV spectrometer that is better than that of SOHO/SUMER (Hinode/EIS) by reducing thenumber of reflections (or reducing the filter components that shut the visible light). A higher
sensitivity by more than a factor of 5 increase from the current space instruments is required. In
addition to measurement of emission-line intensity, a spectral resolution similar to EIS andSUMER is needed to evaluate the dynamics of hot plasmas. The FOV center needs to be changed
by moving an internal optics like SUMER, which could change the FOV in the north-south
direction to investigate the outer region of polar coronal holes. A nominal image size of 256x256
pixel data points with 10 sec cadence produces ~20 kbps data when the data are compressed in a 3
bit/pixel lossy data compression like JPEG2000.

2.1.5 Observables of inner heliospheric imaging (Optional)
The imaging data of the inner heliosphere in the visible-light range are obtained by periodic
exposures that produce a huge data volume. The exposure duration of ~1 min is a typical exposure,and the on-board summation of 60 frames, so that the image producing rate is one imager per 60
min. The final image size after the binning is 1024x1024 with 16 bit depth at each image point.
Assuming 40 percent data compression, the final data producing rate is about 2 kbps. The FOV size
is 60 to cover a large inner heliosphere and the FOV center point to a point above the solar southpole with an offset angle of ~45 (TBD) from the disk center. This setting of the instrument gives
an overhead view of the Sun-Earth space in the ecliptic plane when the spacecraft is in the northern
heliosphere. The spacecraft needs to rotate by 180 around the spacecraft Z axis if the Sun-Earthspace in the ecliptic plane should also be observed when the spacecraft is in the southern
heliosphere.
The following objects are contained in the observed data: solar wind structures and CME (viascattered light by electrons in the heliosphere), diffuse Zodiacal light (via scattered light by dust
particles in the heliosphere), stars, planets of our solar system, comets, and asteroids. The brightest
object is planets as a point source and scattered light by dust particles as a diffuse source. To
determine a dust distribution from the measurement of Zodiacal-light brightness, a special
observation in which the spacecraft Z axis is largely offset to the disk center may be required. Since
it is strongly related to the safety operation of spacecraft, a risky attitude maneuver shall be avoided.

2.1.6 Observables of in-situ measurements (Optional)
Not studied in detail. The measurement of solar wind velocity for main elemental species will bemandatory. The in-situ magnetic field will also be a basic parameter, but the level of ~1 nT
magnetic field at 1 AU may not be measured by the strong magnetic fields from permanent
magnets used in ion engines when the SOLAR-C orbit is formed with help of solar electric
propulsion system.

36

2.1.7 Summary of Required Data
Table 3 summarizes the data required for SOLAR-C Plan-A.
Table 3. Data required for SOLAR-C Plan-A
Measurement

Wavelength(Å)

1Photospheric line-of-sight
Doppler velocity:

Visible(Pixels)

Number ofimagestransmitted

Cadence

512×512

1

1 min

*13

1024 ×1024

1

1 min

*51

1

10 minSpatialsampling

Image Size

4.0 arcsec

Data rate(kbps)

Global Helioseismology

2

Photospheric line-of-sight
Doppler velocity: Local
Helioseismology

Visible3

Photospheric line-of-sightmagnetic field

Visible4

Photospheric vector
magnetic field

1.0 / 2.0
arcsec
1.0 arcsec

2048×2048

2.0 arcsec

1024 ×1024
2048×2048

Visible1.0 arcsec

1.0 arcsec

2048×2048
2048×2048

256×256*10*5

8 hours*7

10 min*5

1 hour

#1 hour

#

5 min

#

1

10 sec*19

16

5

Imaging of Chromosphere

Visible(TBD)

6Imaging of TR and Corona

EUV/Xray~1.2 arcsec

7

Imaging spectroscopy of
Chromosphere, TR, and Corona

EUV~1.0 arcsec

8

Monitoring TSI

-

-

-

-

1 min

2

visible0.03

1024 ×1024

1

1 hour

*1

-

-

-

-

-

~1

9

Heliospheric imaging[option]10 In-situ (TBD)~2.4 arcsec

1024×1024
spatial×spectral
256×256*(#)Data compression down to 3 (1.5) bit/pixel is assumed

37

1

2

3
3

20

2.2 Science Payload

Table 4 shows candidate payloads for SOLAR-C. The description of each payload is shown in thefollowing subsections.
Table 4. SOLAR-C Plan-A Model Payload
HAI (Helioseismic Activity Imager)
Weight65 kg
Aperture: 12.8 cm
Detector: APS or Backside illuminated CCD
Sampling: 0.5 arcsec/pixel (TBD)
Detector format: 4K4K (TBD)
Wavelength: 5000-8700 Å
LIM (Luminosity and Irradiance Monitor)
Weight5.0 kgNumber of cavities
3
EAI (EUV Activity Imager)
Weight
12 kg
Ritchey-Chretien telescope with Mo/Si mirrors
Detector: APS or Backside illuminated CCD
Sampling: 1.2 arcsec/pixel
Detector format: 2K2K (TBD)
Wavebands: 171, 195, 304, 335 (all TBD) Å
ESS (EUV Scanning Spectrometer)
Weight30 kg
Two reflection system: offset parabola + concave grating
Spectrometer with a factor of 4 (TBD) magnification
Detector: intensified APS or intensified CCD detector
Sampling 0.75 arcsec/pixel (TBD)
Detector format: 4K(spectral)1K(spatial) TBD
Waveband (TBD):
IHI (Inner Heliospheric Imager; Optional)
Weight8.0 kgDetector: APS or Backside illuminated CCD
Sampling: 0.03 deg/pixel
Detector format: 2K2K (TBD)
Others (In-situ Measurements etc.)
Weight kg
Mission data processor if required
WeightTotal weight (130 kg tentatively allocated for payload)

MDI+electronics 57 kg
HMI+electronics+ harness 73.5 kg

2 PMOD+electronics in
VIRGO: ~5 kg
TIM(4 cavities): 10 kg
STEREO EUVI
telescope alone: 7.7 kg

Use of MCP
148cm focal length
STEREO HI-1+HI-2+Electronics: 13.8kg

 kg
120++ kg

2.2.1 Helioseismic and Activity Imager
This instrument observes the Sun in the visible light for imaging the Sun in the continuum and atabsorption lines, for acquisition of the helioseismic information from Doppler velocity observations,and for mapping the photospheric magnetic fields. It consists of a refractive telescope, apolarization modulator, a tip-tilt mirror, a narrow band tunable filter, and a mechanical shutter, anda 2D imaging camera. The narrow bandpass filter has a bandpass of 100-200 mÅ in full width at38

half maximum (FWHM), and its central wavelength can be tuned in 5000-8700 (TBD) Å. A similar
type filter with a wider bandpass of 1 Å FWHM was operated for imaging the green-line (5303)
corona for 10 years at a ground-based observatory in Japan.
The baseline detector is a large-format CCD camera with 4K×4K (TBD) pixels. When a scientificgrade CMOS detector is available, it is better to be adopted from the viewpoint of radiation
hardness in the space environment.
The whole-sun area has to be observed in most of the time because the Doppler velocity
observations for helioseismology will mostly be running except for some special operations.
The baseline data recording rate of this instrument is 50 kbps. When one of data is compressed to 3
bits/pixel average, three types of data of 512×512 pixels, magnetograms, Dopplergrams, andcontinuum image, can be obtained every 1 min.
Figure 17 shows the schematic optical diagram of the instrument. One of programmatic difficulties
in developing this instrument is to realize a narrow band tunable Lyot filter that is being developed
to raise TRL. When it is difficult to be developed , the use of Michelson interferometer with awider band Lyot filter like SOHO/MDI and SDO/HMI observing a single absorption line or a
Fabry-Perot filter may become a fallback plan.
The selected line for the Doppler and magnetic field measurement is a TBD line at TBD Å. Ni I
6768 Å (Fe I 6173Å) was selected for SOHO/MDI (SDO/HMI).

Figure 17. Schematic optical diagram of HAI conceptual design.

2.2.2 EUV & X-ray Spectroscopic Imaging Telescopes
A fast cadence imager is needed for tracking the evolution of dynamic events in the outer solar
atmosphere. The SOLAR-C Plan-A is not a mission targeting to resolve elementary structures. It
will be sufficient to cover the whole Sun corona with a range of extended corona above the limb of
the Sun by an imaging system that resolves a few arcsec structures. Observations of solar polar
region from an out-of-ecliptic view are new issues and a telescope to observe the structures near thepolar region is required. A whole Sun imager of ~1 arcsec spatial sampling is proposed here. The
whole Sun with extended corona can be covered by a 2D imaging device of 2K × 2K pixels at the
distance of 1 AU from the Sun. An EUV imaging telescope like STEREO EUVI may be more
39

suitable than the X-ray telescope to observe polar plumes and inter-plume regions in polar coronalholes. In the late 2010’s and early 2020’s whole Sun EUV imagers are planned in Solar Orbiter and
GOES spacecrafts and the presence of these imagers needs to be considered for the selection of thetelescope bandpass. The SOLAR-C Plan-A EUV Activity Imager (Figure 18) is a Ritchey-Chretien
telescope consisting of two hyperboloid mirrors. Each mirror is coated with Mo/Si multi-layers toenhance the reflectance of EUV light. The full width at half maximum of the reflectance as afunction of wavelength is 10–15 Å around a central wavelength of each bandpass. Currently
selected wavelength bands are 171, 195, 304, and 335 Å for their central wavelengths. A
stereoscopic observation with SO will be possible at 171, 304, and 335 Å bands because the same
EUV band with SO is selected. The CMOS APS is a primary candidate for the detector and the fallback plan is a back-illuminated CCD that was used in Hinode, STEREO, and SDO imagers. The
imager on STEREO with a different selection of wavelength band can be a model payload.

Figure 18. (left): SOLAR-C EUV Activity Imager
Figure 19. (right): SOLAR-C EUV Spectrometer

The imaging spectroscopy is a key instrument to detect the source of fast solar wind flows frompolar coronal holes. The following performance is required: (1) the higher-throughput than SOHOSUMER and Hinode EIS for both quiet sun (QS) and active regions (AR), (2) 1 arcsec or slightly
better spatial sampling, (3) a scanning rage of ±4 arcmin in the scanning direction, and (4)simultaneous observations of emission lines formed in the chromosphere, transition region, andcorona. The SOLAR-C Plan-A science payloads have to be set at the side of the Sun, that is thespacecraft +Z direction, to the bus module because the surface of side bus panels located in thespacecraft ±X and ±Y directions are all used as radiators. Hence, the length of this imaging
spectrometer must not be long like Hinode EIS, which is longer than 3.2 m in the Z direction. In
order to realize a short instrument length without adding another reflection, the spectrometer needs
to have a magnification using a single concave grating. The design of this type of spectrometer was
first found by Roger Thomas (Thomas 2003). The instrument, SOLAR-C EUV Spectrometer(Figure 19), consists of a single off-axis parabola as a primary mirror of the telescope, a slit/slot
selector, a toroidal (or ellipsoidal) concave grating with varied line space as a disperser, and solarblind intensified CMOS detectors. Three wavelength bands are chosen to observe thechromosphere, transition region, and corona simultaneously: one is 310-340 Å (TBD), and 760-800
Å (TBD), and 1000-1100 Å (TBD).
The main body of these telescopes and imaging spectrometer are made of CFRP because of theweight reduction. In addition to the stringent contamination control of used materials including
sufficient material bakeout, one point of note is that temperatures of all optical components shouldslightly be higher than that of the main body by the thermal design for a long-term stable
instrument performance.
40

In the imaging spectroscopy with a grating spectrometer, there are two specific directions in theplane of the sky, that is, a slit direction and a scanning direction perpendicular to the slit direction.
The direction along the slit will be in the north-south direction in the imaging spectrometer. It will
be good to rotate the spacecraft around its Z axis for a specific observation to set the slit for any
direction.

2.2.3 Solar Irradiance Monitor
The SOLAR-C total solar irradiance monitor is an absolute radiometer with cone-shape cavities. In
front of each cavity there is a mechanical shutter to control the radiation input from the Sun. Eachcavity has a heater and a thermal sensor, so that each can be a reference cavity. One of cavities isperiodically exposed to the Sun. During the exposure, one of remaining cavities unexposed to the
Sun is heated by the heater attached until the electrical heat input balances the solar input. For thisoperation there needs two cavities in this instrument at least. The cadence of periodic exposure is 1
min (TBD). There is degradation of the cavity surface by being exposing to the space environment,
mainly due to solar UV illumination. To calibrate the degradation of the cavity surface, aninfrequent exposure of a reference cavity to the Sun is needed. Probably, there will be another TSI
instrument in low-Earth orbit at that time, so the SOLAR-C TSI monitor will also be calibrated tothe other monitor when the SOLAR-C spacecraft is located in the ecliptic plane. There are a few
groups to have heritage on developing the space TSI monitor, PMOD, JPL, LASP, and so on. A
cross calibration to a cryogenic absolute radiometer is required before the flight.

2.2.4 Inner Heliospheric Imager (Optional)
A slightly-modified heliospheric imager on STEREO can be an optional candidate payload. The
lens unit consisting of multiple optical lens has a large field of view (FOV) of about 60 and the
FOV center to the Sun is offset by 45 around the spacecraft X-axis looking at space above thesolar south pole when the spacecraft is in the ecliptic plane (Figure 20). When the spacecraft islocated at the northern heliosphere, the heliospheric imager can look at the interplanetary space
between the Sun and Earth. The spacecraft needs to rotate 180 around the spacecraft Z axis when
the same interplanetary space is observed in the southern heliosphere.

Figure 20. Field of View of Inner Heliospheric Imager

The scatted light level is reduced by the multiple baffles located in front of the wide-field lens unit.
A 2D imaging device of 2K × 2K pixels that is cooled down to -60 is placed at the focal plane.
41

The brightest component is from planets in the solar system as a point source and the zodiacal light
as diffused background. The bandpass is 5000-8000 Å (TBD) in the visible range. In order to avoid
the blooming due to a bright object in FOV, a CMOS device will be preferable. Each image is
obtained by ~1 min exposure duration. A multiple set of exposures is summed together by an onboard processing to enhance the component originating from electron scattering. The processeddata are reformed into a telemetry data.
If the FOV of this instrument can point to the opposite side of the ecliptic plane by the spacecraft
maneuver, a basic data set for determining the interplanetary dust distribution may be obtained forthe first time.

2.2.5 In-Situ Instruments (Optional)
The role of in-situ instruments has not been discussed in detail in the SOLAR-C Plan-A working
group. The out-of-ecliptic in-situ measurements were performed in the Ulysses mission. The orbit
inclination is about 80, which mostly covers the whole heliocentric latitude. The orbit is highly
elliptical, so that the there is a period of a fast and a slow latitudinal scan for the heliosphere withdifferent distances from the Sun. One major point of the SOLAR-C Plan-A orbit is that the
spacecraft periodically travels between northern and southern maximum latitudes in a half year atnearly 1AU distance from the Sun. It is possible to do a research on the heliospheric physics thatstrongly depends on the solar cycle. A transport of cosmic rays in a solar cycle may become atarget of SOLAR-C Plan-A regarding the production of cosmogenic isotope like 14C and 10Be,
which are used as indices of the past long-term solar activity.

2.2.6 Other Instruments (Optional)
Since the SOLAR-C Plan-A spacecraft travels in a unique orbit, some instrumental proposals may
arise from non-solar physics community. Although it may be important to consider them here to
maximize a rare opportunity, it is beyond the scope of this proposal.

422.3 Spacecraft Requirements given from the Science Objectives
2.3.1 Attitude Control Requirements
On the SOLAR-C Plan-A there are imaging instruments that require periodic exposures in a long
period of time, so that the spacecraft +Z axis is basically pointed to the Sun by a three-axis
stabilized attitude control system (ACS). Since one of main payloads, which has marginally cover
the whole Sun in the field of view, points to the Sun center for a day, a week, a month or a few
months for the purpose of helioseismic and magnetic field observations, the nominal pointing
direction is the center of the solar disk, so that the pointing coordinates of the other instrumentswith narrower field of view is restricted by the condition. Thus, when the other payload needs tolook at a different location on the Sun outside its narrow field of view, the payload has to have amechanism to change its own field of view. The tracking of the differential rotation of the Sun by
the spacecraft attitude control is not required for observations with a small field of view.
Yohkoh and Hinode also have three-axis stabilized ACS and the spacecraft Y-axis is stabilized topoint to the solar north. The functionality to rotate the spacecraft around the Z axis at any angle isto be requested, and provides a unique opportunity to observe scientific targets for the EUV
scanning spectrometer that has a specific orientation in the plane of the sky. It is the slit direction. If
the inner heliospheric imaging is adopted, the spacecraft rotation of 180 around the Z axis
becomes mandatory and a step-by-step rotation of the spacecraft may sample the scattered photons
from dusts in the solar system.
Table 5. Requirements for Attitude Control Stability (TBD)

Payload- Imager- Spectrometer

Duration

ΔX, ΔY (arcsec)

1 sec

0.6

3σ

10 sec

0.8

3σ

5 min
1 hr

1.6

0-p

24 hr

*2.0

ΔZ (arcsec)

400

Unit

0-p

*: will be better than Hinode due to no external disturbance except for solar radiation pressure in
the interplanetary orbit. An exceptional period happens during the operation unloading theangular momentum of reaction wheels.
The spacecraft pointing stability achieved in Hinode completely satisfies the SOLAR-C Plan-A
requirement. The long-term stability for an order of hour, a few arcsec per hour, will automaticallybe improved in SOLAR-C Plan-A, because the disturbance due to the thermal distortion disappears
in the interplanetary space. Table 5 shows requirements for the attitude control system.

43

2.3.2 Time management
Except for the helioseismic instrument, the ISAS time tag system that was used in Hinode isapplicable to SOLAR-C Plan-A. An additional requirement is given from the helioseismic
instrument that requests a regular interval of exposure sequences. Table 6 summarizes therequirement for the time tag. The information of the time tag is distributed from the spacecraft data
handling unit to the instrument controller of the helioseismic payload.
Table 6. Requirement of Spacecraft Time Tag (TBD)

Short term stability (1 – 100 sec)

± 0.5 ppm TBD

Requirement in SOHOLong term stability (1 year)
Duty cycle± 1.0 ppm TBD
0.5±0.05 sec

Requirement in SOHO
Requirement in SOHO

2.3.3 Data Compression and Data Storage
Since the telemetry should effectively be used, all the data need to be compressed on-board thespacecraft. The 12-bit lossless DPCM and lossy JPEG image compression were adopted in Hinodeand a similar type of image compression is required in SOLAR-C Plan-A mission. It is efficient tocompress the data in one box that is connected to all payloads, but the alternative is that eachpayload produces the compressed data to the spacecraft. The selection of the scheme is TBD.
The spacecraft is accessible from the Earth every day and the downlink will be periodically
performed. But the total telemetry rate at the most distant location is not high and it is the bottle
neck in the high-cadence observations. The following design philosophy will be applicable.
- Each payload can send the data to the spacecraft at any time.
- The data storage area is divided into small number of volume and each payload can send the databy a specified data storage volume.
- When the data are sent to the ground, one of data storage volume is dumped according to aplanned schedule.
Some payload can do a high-cadence observation easily under this system, and the issue is thelargest data size of a data storage volume that is related to the longest duration of the high-cadenceof observation.
When the science data are produced in the rate of 100 kbps at the most distant part in the orbit
during the science operation phase, total data volume per day is 8.2 Gbits (100 kbps × 1024 bit/kbit
× 86400 s = 8.85×109 bits). However, the data rate can be higher when the distance to thespacecraft is shorter or when the Ka-band link condition is more stable than expected. When the
100 kbps data rate is possible at spacecraft position of 40 from the solar equatorial plane, the data
rate can be 2.0 (3.3) times larger at 30 (25). The capacity of the data recorder should be selected
for an observing case of the maximum data rate condition. The tentative number may be 32 Gbits(TBD) including the spacecraft housekeeping data.
44

It will be useful from the operational point of view if the data recorder storage is divided into
several independent volumes or partitions and if each payload can send the data by selecting one ofvolumes.

45

3 Spacecraft System
The spacecraft system design, including the orbit profiles and operation concept, are presented inthis section. Section 3.1 summarizes the requirements to the spacecraft system derived from thediscussion in the previous chapters. Obviously, the most important requirement of SOLAR-C PlanA is to place the spacecraft at a high solar latitude. Working out the way to reach there dominates
the mission design, which is discussed in Section 3.2. Several options are investigated, and an
option using solar electric propulsion (SEP option) is presented as the primary option in this report.
Based on the system requirements and the orbit design results, the spacecraft system for carryingout the mission is briefly investigated. The objective of the study is to confirm the feasibility of themission mainly from the point of view of the resource budget, and its successful result is given in
Section 3.3. The details of the subsystems are to be shown in Section 3.4, At this stage of interim
report, however, the main focus is on those subsystems that require newly developed technology(communication system, light weight solar panel, and electric propulsion system), and the othersubsystems are only briefly touched.

3.1 Spacecraft System Requirements
The major requirements to the spacecraft system (including orbit design) are,
- To reach the solar latitude of 40° by early 2020’s.
- To support science instruments’ mass of 130 kg (or larger).
- To transmit the scientific data continuously recorded in the rate of 100 kbps to the ground.
The third item is interpreted as the 300 kpbs down link rate under the assumption of 8 hours use of
the ground station per day.
The following requirements are also recognized. However, they are regarded as non-dominating in
the total system design and left for the further study in this report.
- Providing payloads a stable platform for imaging observations- 150 W (TBD) payload power- Regular notification of absolute time for regular-interval exposures required by the helioseismic
instrument.

3.2 Orbit and Mission Profile
The SOLAR-C Plan-A is the first Japanese solar mission that operates in the interplanetary space.
The following new issues in mission profile should be noticed when compared with the otherJapanese (Earth orbiting) solar missions.
- It cruises in an interplanetary orbit, and takes much longer time to reach the final orbit.
- It only carries light payloads even if a heavy launcher (as H-IIA) is used.
- It is in severer radiation environments.
The following new issues from other Japanese interplanetary missions should be noticed as well.
- It requires much higher telemetry rate than ever achieved.

46This item is not only reflected in the communication system design, but also considered in the orbit
design to keep the distance from the Earth to the spacecraft short.

3.2.1 Trajectory Options for SOLAR-C Plan-A
The major mission requirement of SOLAR-C Plan-A is to observe the Sun from high latitudes. The
target maximum latitude is tentatively specified as 40° (Figure 21a). To observe the Sun from thehigh latitudes, the space observatory (spacecraft) must be in an orbit largely inclined with respect tothe ecliptic plane. It is not an easy task to inject the spacecraft into the orbit of this type. A roughestimate shows that the velocity increment required to inject the spacecraft into this largely inclined
orbit is approximately 20 km/s (Figure 21b).

v

Figure 21. Schematics of Trajectory Design Challenge in SOLAR-C Plan-A

In order to achieve this severe mission target, possible trajectory sequences are investigatedconsidering the application of various trajectory manipulation techniques. The items considered are,the geometric relation (i.e. the tilt of the solar equatorial plane to the ecliptic plane), launcher
capacity, planetary gravity assists, and the usage of a highly efficient propulsion system. As a result,four major trajectory options are listed up to achieve the mission objective (Figure 22).
The first option is called “SEP option”, which is characterized by the usage of the solar electric
propulsion (SEP) combined with the Earth gravity assists (EGAs). In this option, SEP is used toincrease the relative velocity (v) at the Earth encounter and the direction of v is changed by EGA
to contribute to inclination increase. As a result of the sequential use of this process, latitude 40°
can be achieved in five years from the launch. The other options are characterized by the usage ofthe planetary gravity assists only, without using SEP. The second option is called “Jupiter option”.
In this option the operation sequence begins by using the Jupiter gravity assist (JGA) to highly
incline the orbit plane from the ecliptic plane. What is unique to this option is the usage of EGAs inorder to reduce the orbit period after JGA, which enables the frequent observation of the solar polar
region. The remaining two options are characterized by the usage of EGAs and the Venus gravity
assists (VGAs). In both options the operation sequence begins with the sequence of EGAs/VGAs to
increase v to the planets. Then, the sequential planetary gravity assists are used to change thedirection of v so as to contribute to incline the orbit. “Venus-1 Option” uses the Venus for theplane change gravity assists, whereas “Venus-2 Option” uses the Earth for this purpose.
47

Among the four options, the first “SEP option” has a major advantage in complying with the mostof the mission requirements (the detail shown in 3.1.2). The option has frequent launch windows(every half year), which should be also counted as a merit from the programmatic point of view. On
the other hand, it must be noted that this option requires an advanced technology, which is not
required in the other options. This point will be discussed in 3.3. In view of the major advantage
mentioned above, in the following part of Section 3, the case of SEP option is focused on.
Alternative Jupiter option with the same final orbit is shown in Appendix B.

Figure 22. Trajectory Options for SOLAR-C Plan-A

3.2.2 Orbit and Mission Profile of SEP Option
Before showing the baseline sequence of the SEP option, a couple of its variants are introduced
briefly. The first method, which is called “Direct Inclining Method (DIM)”, is a simple one whichuses SEP directly to increase the inclination. The second method uses SEP combined with EGA,
which is called Electric Propulsion Delta-V Earth Gravity Assist (EDVEGA). The method uses48

EDVEGA repetitively, and is called “Sequential EDVEGA Method (SEM)”. Both methods were
investigated quantitatively, and it was concluded that, DIM is infeasible from the points of view ofthe spacecraft’s mass budget and the operation time of the ion engine system (IES), whereas SEM
is feasible in these aspects. A slight variation of SEM is also investigated, which applies anadditional Venus gravity assist (VGA) prior to the sequential EDVEGA. This method has anadvantage in reducing the launch energy drastically compared to the original SEM. However, theuse of Venus/Earth gravity assists limits the launch opportunity, which makes it difficult to takeadvantage of the geometrical relation (i.e. the tilt of the solar equatorial plane to the ecliptic plane).
From this point, the usage of VGA is regarded only as a backup option.
As a result of the discussion above, SEM is adopted and used to construct the baseline sequence of
the SEP option. The basic procedure of SEM is described as follows.
1. The spacecraft is injected into the Earth synchronous orbit to re-encounter the Earth after one
year cruise.
2. During the cruise, SEP is used to maximize the spacecraft’s relative velocity to the Earth (v) at
the next Earth encounter. Note that the thrust does not necessarily increase the inclination by
itself. To enhance the efficiency to increase v, an elliptic orbit is used for the cruise orbit.
3. By EGA, the direction of v is changed to contribute to the inclination increase.
4. By the repetitive use of the steps 2 and 3, the inclination is increased step by step.
Before presenting the constructed trajectory sequence, major assumptions used in the trajectory
design should be noted.
The first assumption is related to the launch condition. The assumed launcher is the Japanese H-II
A heavy launch vehicle equipped with a solid motor upper stage. The initial mass of the spacecraft
is assumed to be 1200 kg. Assuming the practical settings of the launch site and the launch
direction, the launcher is capable of injecting the spacecraft into the escape orbit with the excessive
velocity (v) of 7.3 km/s. The launch date is selected to take advantage of the not negligible tilt
(7°.25) of the solar equatorial plane to the ecliptic plane, and they are June 7 or December 8. In thefollowing discussions, June 7 is used as the launch date.
The second assumption is related to the ion engine system (IES). The specific impulse (Isp) of IES
is assumed to be 3800 s, and the maximum thrust of IES is assumed to be 120mN in total. In the
trajectory design, the actual thrust available for the maneuver is constrained by the available power,
which is assumed to decrease as the inverse square of the distance from the Sun.

49

Figure 23. Trajectory profile of SEP option

The third assumption is related to the eccentricity (e) of the cruise orbit. To efficiently increase v,the eccentricity e should be large. However, considering the difficulties in the thermal design of thespacecraft, the eccentricity e is constrained to be less than 0.3 in the trajectory design.
Figure 23 shows the trajectory profile of the SEP option. In Figure 23a, the trajectory of the first
revolution is projected on the ecliptic plane. The figure shows that the spacecraft’s orbit has
eccentricity, and intersects with the ecliptic approximately at the Earth’s position at the launch. The
acceleration vectors indicate that the thrust is used to decelerate the spacecraft at aphelion and toaccelerate the spacecraft at perihelion, which result in the increase of v at the next Earth encounter.
Figure 23b shows the trajectory profile through the sequence. The trajectory is projected on theplane perpendicular to the ascending node direction so that the gradual change of the orbit plane
can be easily seen. It is observed that the first four cycles have asymmetry resulted from the orbit
eccentricity. However, the orbit is finally circularized by EGAs, which results in the symmetry
observed in the final orbit (the orbit whose inclination is the highest).
Table 7. Sequence of events of SEP option

50

The sequence of events of the SEP option is summarized in Table 7. In the table, v denotes therelative velocity to the Earth at EGA, and iSEQ denotes the inclination to the solar equatorial plane.
The spacecraft reaches iSEQ of 30° after the 3rd EGA (3 years from the launch), and finally reaches
iSEQ of 40° after the 5th EGA (5 years from the launch).
Figure 24 shows the profiles of some important parameters of the mission. The top chart indicatesthe points of events and basic operation concept. The following three charts respectively show the
profile of the spacecraft’s distance from the Sun and the Earth, its instantaneous solar latitude, andthe expected down link rate of the scientific data.

cFigure 24. Mission profile of SEP option

In the first three years from the launch, before the spacecraft reaches iSEQ of 30°, the spacecraft
operation is devoted to increase the inclination. IES is operating most of the time, and silent
condition for the scientific observation is basically not guaranteed in this phase. Even in this phase,intermittent suspension of IES is planned for accurate orbit determination, which may be used as51

occasions for observations. In the fourth year, after EGA #3, iSEQ exceeds 30°. From this year
onwards, the period during which the spacecraft is at latitudes higher than 30° is allocated as “theobservation phase (orange area in Figure 24)”, and the IES operation is intentionally suspended.
Even after EGA #3, IES operation is continued while the spacecraft is in the lower solar latitude.
This IES operation and the following two EGAs contribute to the further increase of the inclinationand the circularization of the orbit. Finally, as a result of EGA #5, the spacecraft is injected in to thefinal observing orbit, the circular one-year orbit with iSEQ of 40°.

3.3 Spacecraft System Design
The SOLAR-C Plan-A is the first Japanese solar mission that operates in the interplanetary space.
The following new issues in spacecraft system design should be noticed when compared with theother Japanese (Earth orbiting) solar missions (a part of issues results from the mission design).
- It requires a much longer mission life.
- It is in severer radiation environments.
- It carries light payloads only.
- It requires large solar array panels- It requires fuel to keep three-axis attitude control.
The following new issue from other Japanese interplanetary missions should be noticed as well.
- It requires much higher telemetry rate than ever achieved.
This item is already considered in the orbit design, and is also reflected on the communication
system design. In addition, it should be noted that, as a result of using a liquid fuel rocket (H-IIA),
the mechanical environment at the launch is relaxed from that in case of using solid rocket as
previous solar missions.

3.3.1 Spacecraft Configuration
The mission design provides basic information of the spacecraft configuration. The required thrust
determines the configuration of IES, and the configuration of IES determines the required power
which finally determines the size of solar array panels (SAP). The orbit design provides thegeometrical relation around the spacecraft during the mission, which determines the basic locationof the key components. As a result, the basic configuration the spacecraft is determined as shown in
Figure 25.

52

Figure 25. Spacecraft configuration of SEP option

The spacecraft’s +Z-direction is always pointing to the Sun. The scientific payload assembly ismounted on the +Z panel of the bus module and pointing to the Sun. SAP are deployed from thebus module to the ±Y-direction and fixed to the body, facing the +Z-direction. In the trajectory
design process, IES thrust direction is constrained to be approximately normal to the Sun direction,
which enables IES module to be mounted on the –X panel. The panels not facing to the Sun
direction (±Y, –X, -Z) are used as radiator panels to dissipate the large heat generated by IES
components. The geometrical relation among the spacecraft, the Sun and the Earth varies as afunction of time. Therefore, a high-gain antenna for the high rate data transmission must be
maneuverable, and mounted on the +X-panel.

3.3.2 Mass Budget
In order to confirm the feasibility of the mission, the spacecraft mass budget is estimated. In theinvestigation, the total mass of the spacecraft is assumed to be 1,200 kg (which complies with theassumption of the trajectory design), and the difference with the sum of subsystem masses iscounted as margin.
For most of the subsystems, the composition of the components is considered, and the mass ofcomponents is summed up to the subsystem mass, while the mass of subsystems such as structure,harness, etc. is estimated parametrically based on the previous spacecrafts. The propellant mass for
IES is derived from the trajectory design result.
The resulting mass budget is summarized in Table 8. The positive margin of 106 kg indicates thefeasibility of the mission from the point of the spacecraft’s mass budget.

53

Table 8. Spacecraft mass budget of SEP option

3.3.3 Power Budget
In order to estimate the size and mass of SAP, the required power for the spacecraft is estimated. In
the investigation, two cases of power budget (IES operating / non-operating) are estimated. The
composition of the components is considered, and the powers for the components are summed up to
the subsystem power. The resulting power budget is summarized in Table 9.
Table 9. Spacecraft power budget of SEP option(Heater powers for payload not included.)

The mass of SAP is estimated on the basis of supplying the power for IES operating case at 1 AU.
The use of newly developed light weight SAP is assumed, and the specific value of 100W/kg is54

used in the estimation. The resulting mass of SAP is included in the power supply system in themass budget.

3.3.4 Communication Link Budget
In order to confirm the feasibility of the high rate data link, the communication link budget isestimated under the assumed communication system configuration. In the investigation, two casesof link budget (X-band / Ka-band) are estimated. The maximum distance of 0.56 AU in the final
orbit is used, and NASA/DSN 34m antenna is assumed as the ground station. The resulting linkbudget is summarized in Table 10.
Table 10. Communication link budget of SEP option

Item
Frequency
Transmitter power
Transmit antenna gain
Communication distance
Data rate
Link margin

X-band
8400
40.0
36.5
0.56
300k
1.0

Ka-band
32300
20.0
48.2
0.56
1M
2.0

unit
MHz
W
dBi
AU
bps
dB

Received C/N0
Required C/N0

59.3
58.4

64.6
62.6

dBHz
dBHz

The positive margin of 1 dB for 300 kbps downlink indicates the feasibility of satisfying the system
requirement in the final orbit. It should be noted that this downlink rate is achieved by more reliable
X-band link. In addition, 1 Mbps downlink by Ka-band link is expected with positive margin of 2
dB.

3.3.5 System Thermal Design
In the cruise phase of SEP option, the minimum solar distance (i.e. perihelion distance) is to be 0.7
AU. It is assumed that three μ20 ion engines are simultaneously used even at the perihelion passage.
The amount of heat dissipated from IES is estimated to be approximately 1,800 W in total, and it is
not easy to radiate this amount of heat to space. In order to confirm the feasibility of the thermal
design, system thermal analysis is conducted. In the analysis, the simultaneous use of the three μ20
ion engines at 0.7 AU is assumed as the worst case. The Sun-light is assumed to input from thedirection 5° apart from +Z direction, which causes the heat input from the Sun into the side panelsof spacecraft.
The thermal mathematical model consisting of the spacecraft bus module and ion engine module isconstructed, and a steady state thermal analysis is conducted. The result shows that we can still55

have a feasible solution under this worst condition. The solution requires that concentrated heatdissipation from the ion-engine related components to be averaged at each structure panel by heat
pipes and thermal doublers.

3.4 Spacecraft Subsystems
The spacecraft subsystems are discussed in this section. At this stage of the interim report, theinvestigation gives priority to the subsystems that require new key technologies to achieve thismission. They are Ka/X communication system, ultra-lightweight solar panels, and μ20 ion engine
system. The details of these subsystems are given in Sections 3.4.2, 3.4.3 and 3.4.6 respectively. As
for the other subsystems, the functions required to them are supposed to be achieved with proventechnologies. Their compositions are tentatively assumed based on the previous missions, and they
are only briefly touched in this report.

3.4.1 Data Handling System
The data handling subsystem are supposed to be composed based on the heritage of the previous
interplanetary spacecrafts, Hayabusa or Akatsuki. It is composed of DHU (data handling unit) and
TCIM (telemetry command interface module). DHU manages the automatic/autonomous functionof the spacecraft bus system, collects and records the HK (house keeping) data, and manages DH
network as the host. TCIM interfaces telemetry and command between DHU and thecommunication system. It generates the transfer frame and multiplexes the telemetry packets, andoutput to the transmitter. On the other hand, it modulates the uplink signal input from responder,and output command packets to DHU.

3.4.2 Communication System
Since the spacecraft is in the interplanetary space for the out-of-ecliptic mission, a high-speed
telemetry is mandatory for imaging observations of the Sun. We have tentatively set a straw-man
data recording rate of 100 kbps in X-band considering the continuous helioseismic observations.
This level has been achieved in the NASA STEREO mission in the distance of 0.5 AU from the
Earth. On the other hand, the Japanese interplanetary missions that have so far been flown only
achieved the telemetry rate of less than 10 kbps using 64m-aperture downlink station at Usuda deep
space center. When the downlink duration is limited to 8 hours per day for example, 300 kbps
telemetry rate is required for 100 kbps recording rate. If this should be achieved by a domestic
effort, a development of a high-power transponder will be required in addition to reducing any
transmission losses.
Based on the consideration above, the following design guideline is derived.
- Achieve 300 kpbs downlink rate assuming the downlink duration of 8 hours per day.
- Use X-band system heritage in interplanetary communication for stable and assured downlink.
- Use Ka-band system for a higher-telemetry capability under good atmospheric conditions.
56

- Adopt domestic supplied components which we can have by around year 2012.
- Adopt a high gain antenna of 1.6 m diameter.
- Adopt a high-power (40W) X-band transmitter as a natural extension from the JAXA heritage.
- Reduce feeder loss compared with the previous JAXA interplanetary missions.
- Improve the coding gain by the modification of coding method.
The block diagram of the communication system (with the heritages) is shown in Figure 26.
In order to have a higher data transmission rate or a shorter downlink duration, the use of Ka-band
system is supposed in addition to the X-band system. The downlink of Ka-band telemetry is
possible at the NASA DSN stations to track the STEREO spacecrafts. Since the availability of Kaband is affected by the atmospheric condition on the Earth, the Ka-band alone cannot be a solution
to the requirement set by the continuous helioseismic observations.

KaTWT
KaEPC

X/Ka HGAKaSW-1

HGA Control Unit
X/Ka Feeder

XEPC-1

XTWT-1
HYB

XSW-4
DIP-1
X LGA-1

XEPC-2

XTWT-2

XTRP-1

XSW-3

XSW-2

KaTX-1DIP-2

KaTX-2
XSW-5
XTRP-2XSW-1

XLGA-2

Figure 26. SOLAR-C Communication System

3.4.3 Power Supply System
Power supply system of SOLAR-C Plan-A consists of Power Control Unit (PCU), Series Switching
Regulator (SSR), Solar Array Paddle (SAP), Battery (BAT), and Pyrotechnic Controller (PYC). In
the system, the output voltage of SAP is lowered and stabilized by SSR. The functions of PCU are
SSR control, BAT charge/discharge control and management, and the bus power distribution. BAT
is used for power supply when SAP cannot supply sufficient power. The cases supposed are, thelaunch phase, the attitude anomaly, and the large attitude maneuver. The use of Lithium ion battery
is supposed based on the heritage of the previous interplanetary spacecrafts, Hayabusa or Akatsuki.
SOLAR-C Plan-A spacecraft is propelled by the ion engine system. In order to operate three μ20
ion engines simultaneously, approximately 6 kW of power is required to be supplied by SAP.
When conventional rigid type SAP is adopted, its mass is estimated to exceed 100 kg. On the other57

hand, the reduction of components’ mass is strongly requested as a general nature of interplanetary
spacecraft. From this consideration, the newly developed ultra-lightweight solar panel is used tomeet the requirements of high specific power (W/kg) and low storage volume. This component is
identified as the key technologies to achieve this mission.
The ultra-lightweight solar panel is now under development in ARD/JAXA, and uses newly
developed space solar sheet supported by a frame-type structure (Figure 27). It is targeting thespecific power greater than 100W/kg in panel level. Future perspective of the ultra-light weightsolar panel development is shown in Figure 28. Breadboard model was developed last year, and aseries of development tests (vibration, thermal vacuum, and deployment test) were successfully
conducted. The technical readiness level (TRL) is estimated to be 6+α.

Specific Power, W/kg

Figure 27. Ultra-Lightweight Solar Panel

150
Present state
100W/kg~150W/kg

100>120W/kg

Standard Panel
1m  1m unit

BBM (small sized panel)

50

Panel
Structure

Cell

88W/kg-Framestructure-Framestructure-Lateralpanels-Framestructure-Lateralpanels-Framestructure-Lateralpanels-Pantographic
mechanism

2J

2J

3J

3J

Figure 28. Future Prospects of Ultra-Lightweight Solar Panel Development

As mentioned in Section 3.3.1, in the configuration of this spacecraft, Z-axis (the direction
perpendicular to SAP) is basically pointed toward the Sun. Therefore, SAP is fixed to thespacecraft body, and no driving mechanisms are installed.
Li-ion battery is adopted as a high-density power cell in the Plan-A spacecraft. The overcharge andover-discharge protections are dealt by the Series Switching Regulator (SSR) that has beendeveloped in the JAXA MMO mission. The bus voltage regulation is made by SSR monitoring theelectric current at BAT. The period in which the spacecraft is operated by the power from BAT is58

(1) during initial phase just after launch, (2) during orbit control without using the ion-engine
system, and (3) during the safety state in an attitude control for emergency cases. The eclipse
period in which the Sun cannot be seen does not happen in the SOLAR-C Plan-A orbit except for
the launch operation. The capacity of battery is selected by these conditions and it is found to beabout 27 Ah (TBC) in the first estimate.

3.4.4 Attitude Control System
The requirements for the attitude control are summarized in Table 5. The three-axis stabilization isrequired for imaging observations of an arcsec spatial resolution. The attitude control system that is
tentatively assumed here is what was used in Hinode mission. Since the short-term performance ofthe Hinode attitude control system completely satisfies the SOLAR-C Plan-A requirements, no
improvement is necessary. The external disturbance in the interplanetary space is much smaller andmore stable than that in the low-Earth orbit, so that the main source of disturbance is themechanical motions within the spacecraft. One of the differences between SOLAR-C Plan-A and
the Japanese solar missions in low-Earth orbits is the use of chemical thrusters to unload theangular momentum accumulated in the reaction wheels (RW). According to a simple calculation,
the frequency of the RW unloading is estimated to be once per day in a nominal science operation.
The helioseismic instrument may need to have its own image stabilization tip-tilt mirror tocompensate the long-term pointing change or to stabilize the unexpected low-frequency disturbance.

3.4.5 Chemical Propulsion System
Chemical propulsion subsystem is used for unloading of the RW accumulated angular momentum
and the attitude maneuver which requires large control torque (such as initial attitude establishment
and safe hold mode operation). The thrusters are also used for the short term translational motion
control such as the trajectory correction before the Earth gravity assists. Considering the level ofthe total impulse required, the conventional monopropellant (N2H4) system is assumed at thepresent phase.

3.4.6 Ion Engine System
Owing to the success of Japanese asteroid explorer HAYABUSA, which has an ion-engine μ10
system for interplanetary cruise, ISAS/JAXA has a heritage of ion engines. ISAS is developing thescale-up version of the ion engine thruster, μ20. Figure 29 shows the engineering-model thruster,which is in an endurance running test under a vacuum condition. We assume a SEP system inwhich four sets of ion engines are set on the spacecraft and three out of four can be turned on
simultaneously. In the orbit design we assume the thrust of 40 mN and specific impulse of 3,800 s
for each thruster or the maximum thrust of 120 mN for simultaneous use of three thrusters.

59

Figure 29. Engineering model of μ20 ion engine thruster

The μ20 ion engine system consists of Ion Thruster Control Unit (ITCU), IES Power Unit (IPU),
Propellant Management Unit (PMU), Ion Thruster Assembly (ITA), Microwave Supply Unit
(MSU), and IES Pointing Mechanism (IPM). Xe is the propellant of this engine. Xe ions areproduced in a process of discharge by microwaves and are accelerated by electric fields between
two grids. The role of each unit is summarized as follows:
IPM: This is a pointing mechanism of ITA.
IPU: This produces a high voltage (~2.3 kV) to accelerate Xe ions.
ITA: This unit produces thrust from Xe propellant, microwaves, and high-voltage applied.
ITCU: This unit controls the whole ion-engine system (IES).
MSU: This unit supplies microwaves to ionize Xe for acceleration and to make electrons forneutralization.
PMU: This unit provides Xe propellant to ITA.
Each ion engine thruster produces the thrust of 40 mN. By simultaneous use of three μ20 ion
engines, one of which is currently being developed and tested at ISAS, the maximum thrust of 120
mN is obtained. Some components are set on the spacecraft bus structure, and some are set on the
Ion Engine Module (IEM) that is an independent box located outside the spacecraft bus structure.
IEM has six ion engine thrusters, two of which are redundant engines for unexpected situations.
The IPM connects IEM with the spacecraft, and it also acts as a pointing mechanism that slightly
tilts the whole IEM to control the direction of thrust under combinations of multiple ion engine
thrusters, each has a slightly different thrust and is located at a different distance from the center ofgravity of the spacecraft.
Power Control Unit (PCU) on the spacecraft provides the power of 350 W (1,520 W) with MSU(IPU) for operating each ion engine thruster. 1,870 W electric power is required for operating asingle ion engine thruster. It was 350 W for Hayabusa’s μ10 ion engine system. Larger power
consumption in the SOLAR-C Plan-A system gives large heat generation that needs to be radiatedaway from the spacecraft. The ion engine system in the Plan-A spacecraft produces a huge amountof heat that cannot be radiated away from the surface of IEM alone. In order to have a solution for60

the heat dissipation problem, some components, ITA and MSU, are set in IEM, and othercomponents, ITCU, IPU, and PMU, are mounted in the spacecraft bus structure.
In the following points the system is different from Hayabusa’s ion-engine system. These are all
related to the increase in weight or heat dissipation to achieve high thrust in good propellant
efficiency.
1. Size of ITA from 10 cm diameter in Hayabusa to 20 cm diameter in SOLAR-C Plan-A
2. Operating power for ion acceleration
3. Pointing mechanism at IPM
Currently estimated specifications summarized in Table 11 assume higher specific impulse options
that will soon be applied to the EM thruster under endurance test. Original test conditions of the
EM thruster's ion source were 1,300 V (Screen Voltage), 30 mN (Thrust) and 2,800 sec (Specific
Impulse), respectively, and it accumulated 10,000 hours of operation by Sep. 2010. The technicalreadiness level (TRL) is estimated to be 4-5.

Table 11. SOLAR-C Plan-A Ion Engine System

Propellant
Specific impulse
ThrustScreen Voltage
Number of thrusters
Maximum number of thrusters
that run simultaneously
Size of Ion-Engine Module
Pointing mechanism
WeightWeight of Fuel (Xe)
Power

Xe
3800 s
40 mN/engine
2300 V
4
3
1,2001,500600 mm
Newly proposed linear actuators
Range: 5
162 kg
236 kg5,650 W
Regulated bus:
1,090 W
Non-regulated bus: 4,560 W

Need a development
Estimate in 2010
Estimate in 2010
Estimate in 2010

3.4.7 Structure System
A box-type panel structure is adopted as the spacecraft main body (bus module). The approximate
size of the main body is 2m×1.2m×1.5m. The scientific payload assembly is mounted on the +Z
panel of the bus module, and the rocket attachment ring is installed on the –Z panel. The movable
IES module is installed on the –X panel. Two axes pointing mechanism enables to direct the IES
thrust axis to the spacecraft center of mass. The major deployable appendages are two wings of61solar array panels and Ka/X-band high gain antenna. The solar array panels are installed on the ±Y
panels, and they are latched and fixed to the body after deployment. Two axes gimbaled movablehigh gain antenna is installed on the +X panel, and is deployed after the launch.

3.4.8 Thermal Control System
SOLAR-C Plan-A spacecraft is propelled by the ion engine system. When three μ20 ion engines are
simultaneously operated, the heat dissipated from IES is estimated as approximately 1,800W in
total. The major theme for the thermal control system is to radiate this large amount of heat tospace. In order to achieve this objective, large areas on the structure panels that are not facing with
the Sun (±Y, –X, -Z) are used as radiator panels. Concentrated heat dissipation from IES istransferred and averaged at each structure panel by heat pipes and thermal doublers. The area
outside the radiators on the structure panels is covered with the multi layer insulator (MLI), tointerrupt unnecessary heat input/output.
When not all the ion engines are operated, the reduced heat dissipation from IES is compensated byheaters. The heater control electronics (HCE) automatically controls the duty of the heaters basedon their priority. It enables the appropriate management of the power allocation between thermal
control and IES operation even when the available power is limited around the aphelion.

62

Appendix
Appendix A:
A.1 Exploration of α-effect and turbulent diffusion
While it is generally accepted that differential rotation is the main player transforming poloidal into
toroidal field the details of the processes rebuilding the poloidal field from toroidal field are still
uncertain. In the mean-field language these processes are formally described as α-effect. In the nearsurface layers it is in principle possible to observationally constrain the amplitude of turbulent flux
dispersal from models that describe the surface evolution of magnetic field.
In the context of solar dynamo models the following mechanisms for the α-effect are generally
considered: 1. Helical turbulence, 2. MHD shear flow instabilities in the tachocline, 3. Rising flux
tubes in the convection zone influenced by Coriolis force and convection (Babcock-Leighton).
While all these processes are likely to contribute, their amplitude and spatial distribution is notknown well enough to clearly quantify their individual role. Furthermore, recent research points
toward a highly non-linear and also time dependent α-effect resulting from additional constraints
due to conservation of magnetic helicity.
In general it is not possible to verify or rule out α-effect mechanisms within the convection zone
through photospheric observations. However, some insight into the poloidal field generation
mechanism might be possible. In particular measurements of the transport of magnetic flux in thepolar regions by super-granular flows and meridional circulations will provide insight into
efficiency of the Babcock-Leighton mechanism. Note that there is a strong correlation between theway the surface Babcock-Leighton term is modeled and quenched (in 2-D mean field models) andthe resulting amplitude of the poloidal field near the poles. Precise observations of the timing,
structure and strength of the poloidal field is thus key to constrain this fundamental ingredient insolar dynamo models.

A.2 Flows associated with flux emergence
An integral part of the current paradigm is the flux emergence process which consists of thebuoyant rise of strong field from the base of the convection zone toward the solar surface andultimately the formation of active regions.
Direct observations to confirm or disprove this picture are still rare but are of fundamental
importance and need to be improved systematically with SOLAR-C In addition to jet like flows ofa few 100 m/s that might be present at the base of the convection zone, a rising flux tube in theconvective envelope induces additional motions with speeds in excess of 100 m/s based on thintube simulations as well as more recent 3D simulations. This flow is a superposition of theemerging upward motion and the downdraft due to the gravity and has slightly asymmetricstructure in the east-west directions. A flux tube is expected to spend most of the time (about 1
month) in the lower part of the convection zone and only a few days in the upper most 20 Mm.
63However, the rise speed in the near surface layers can reach values of about 1 km/s and is therefore
more likely detectable.
Helioseismic signatures of sunspots have a large contribution from the near surface regions, where
the relative perturbations of wave speeds are maximal. Our ability to seismically image thesubsurface structure of sunspots depends strongly on a good understanding of near surface effects.
The primary source of near surface thermal perturbations is caused by the Wilson depression.
Independent measurements of this contribution in the near surface layers allows us to refine
helioseismic tools and improve their focus on magnetic perturbations in the deeper layers. The
second vantage point of a high latitude mission would allow for stereoscopic measurements of the
Wilson depression, provided the imaging quality of the telescope is sufficient to resolve umbral dot
like features in a sunspot umbra.
Helical turbulence generates helical mean fields. Since magnetic helicity is conserved in highly
conductive fluids, small-scale magnetic helicity of the opposite sign must be shed through thephotosphere or dissipated through Ohmic diffusion in order to sustain the large-scale dynamo.
Evidence for helical field topologies or small-scale magnetic helicity flux through the photospheremight therefore be construed as an observational signature of the turbulent α-effect.
Vector magnetograms will also help quantify the global distribution and flux of magnetic helicity inand through the photosphere which has profound implications for the operation of the turbulent
dynamo (Section3.3). High-latitude magnetograms will also enable a more thorough evaluation of
the effects of surface magnetism on global oscillation frequencies, potentially improving thesensitivity of global rotational inversions. This is crucial for detecting the subtle signatures of zonal
flows possibly associated with subsurface toroidal magnetic flux. In addition, observations of polar
magnetism and related emission on the Sun will help interpret stellar observations by calibratingproxies used to infer magnetic activity.
Magnetic helicity is thought to play a key role on the operation of small versus large scale dynamo
in the Sun. It can be demonstrated that for fast processes with respect to magnetic diffusion and bydefining a reference magnetic field, a gauge invariant magnetic helicity can be defined in diffusive
open system such as the Sun. It is also well known that stability conditions of the surface magnetic
structures are dependent on their magnetic helicity budget and the injection of magnetic flux andhelicity by the convection. This quantity can thus serve a very important role since it can relate
inner and outer magnetism helping us to draw and improve our theoretical view of the solar
magnetism in a more integrated way.
While it is in principle possible to detect helical flows in the convection zone or nonaxisymmetricflows resulting from MHD instabilities in the tachocline, it is not clear how to relate these flows toan α-effect without making very strong assumptions (first order smoothing, second ordercorrelation approximation), which cannot be justified for the solar parameter range from firstprinciples. In that sense a measurement of these flows can be taken only as an indication for apotential α-effect in these regions. On the other hand the contribution from the Babcock-Leighton
α-effect can be constrained by observations of the flux-budget in the photosphere. Pure BabcockLeighton flux-transport dynamo models typically lead to a polar (mean) field strength of about 100
Gauss (if solutions are normalized to about 104 Gauss at the base of the convection zone) clearly inexcess of the observational constraints around 10 Gauss. The latter shows that likely additional αeffects within the convection zone or tachocline exist and it has been shown that their inclusion
64

leads to models consistent with surface observations. A detailed study of the surface flux budget at
all latitudes will promote a better understanding of the mechanisms of poloidal field regeneration in
the sun. In particular, a quantitative determination of the contribution from Babcock-Leighton
sources will impose additional constraints on (unobservable) sources within the convection zone
and tachocline. Meanwhile, observations of magnetic helicity may be used to help assess the role ofthe turbulent α−effect (Section 3.3).

A.3 The inclination requirement

Figure 30. Travel-time perturbation 2δτ computed for various combination of position and thickness ofsoundspeed perturbation of the size 10-5, against full skip angle 2Φ.

We have so far established that observing the sun from a high-inclination orbit may well be
desirable in our attempt to study dynamical processes in the solar interior and ultimately understandthe solar dynamo. The advantage we will have, however, would not be by orders of magnitude interms of our primary observables such as, say, travel-time measurement, because projection effectis just a matter of a factor except near the very edge, where Doppler signal would be dominated byhorizontal motion brought about by convection. Let us look at an example.
According to the ray theory, a ray that reaches tachocline, around r/Rsun~0.7 has a skip angle of~45. If one is aiming at observing tachocline at the solar equator and immediately beneath the disccentre, the product of two foreshortening factors, by which we lose special resolution and thevertical component of velocity, is only cos2(22.5) ~ 0.85. It is generally believed, however, that atleast some of the key dynamical features are at higher latitudes. For a ray that begins at 7.5latitude, travels down to tachocline at 30 latitude and returns to the surface at latitude 52.5, during
all which staying within the meridian, the product is cos (7.5) cos2(52.5) ~ 0.6. For a feature athigher latitudes, the figure will be worse but, it is clear that we are talking only about improvement
by a factor 2 or 3. We would like to point out, however, that even a factor of 2 may turn out to be
crucial at the end of the line of our investigation into the deep interior.
65

More complete analysis of detection capability for a certain signature in the solar interior, by timedistance helioseismology, should take the following steps.
1. Travel-time perturbation: calculate the magnitude of travel-time perturbation for a given pair ofpoints on the solar surface for which cross-correlation function is computed. This can be done by asimple ray approximation, by a more sophisticated method such as Born approximation, or bydirect numerical simulation of wave propagation.
An example-- travel-time perturbation due to sound speed anomaly near the base of the solar
convection zone (Figure 30). Perturbation to travel time 2τ (here τ is the travel time from thesurface to the inner turning points) has been computed for the relative sound speed perturbationΔc/c = 10−5, and various position and thickness, by ray approximation. The ray path Γ isdetermined by local group velocity dω/dk of pure acoustic waves for which local dispersion
relation ω=kc holds. Then soundspeed anomaly Δc/c is integrated over the path Γ. For flow of asimilar size, measured in v/c, the integration has to be done for v cos μ/c, where μ is the angle
between the path and the flow velocity. There is the reduction factor |cos μ | but if integrated overtheta (see Averaging below), by rotating ray path, and if we optimize the signal, the integrated
reduction factor would be just 2/π.
2. Uncertainty: calculate the uncertainty in travel time obtained for a given pair. If the pair is close
to the disc centre, it is known that the so-called realization noise, caused by stochastic excitation, isthe dominant source of contribution. Closer to the limb, however, and particularly if the oscillation
measurement is done in Doppler, granulation noise can be important. Other quantities that are ofimportance to uncertainty are spatial resolution, cadence and length of observations.
The way a poor spatial resolution degrade the travel-time measurement, for example, is not througha straightly enhanced statistical uncertainties, which in fact would generally be decreased. When agroup of travelling wave packets is measured with a poor resolution, there would be many
components travelling between a given pair of pixels, with widely different travel time. As thedispersion in the travel time increase, the travel time between the pixels becomes increasingly illdefined and therefore difficult to measure.
3. Averaging: Most of the time, a single measurement for just a pair of point is unlikely to yield
sufficient signal-to-noise ratio. The averaging may be done in space, or in direction. To account of
the effect of averaging we need to know error correlation between quantities that are (probably
weighted-)summed.
A note on the direction average -- this is averaging the signal over a ray path that is rotated aroundthe target point. Part of the averaging is about averaging product of projection factor for the pair ofsurface points where the individual ray begins and ends. As is seen in Figure 31, for annuli of themoderate size, it is adequate to approximate this by squaring the projection factor for the mid point(the surface point directly above the target), or a certainly point close by.
On the other hand, one can produce a fairly robust argument for the inclination requirement bylimiting the objective to detection of meridional flow. The meridional flow velocity below 60
degree latitude has been measured by SOHO/MDI. From 30 latitude, therefore, one would be able
to measure meridional flow up to the pole, one might think. It is true that from 30 latitude we66

observe a surface element at the pole as well as we would a surface element at 60 latitude from theecliptic plane. As has already been mentioned in the above, however, there is an issue of averaging.
The result obtained by Giles et al (1997) was produced by a longitudinal average within 50 fromthe central meridian. With the increasing latitude μ, the amount of averaging one can carry out
would decrease as cos μ even if we keep the longitudinal resolution, which we will not. Let us say
we would like to do at 80 latitude as well as SOHO/MDI did at the 60 latitude. By noting(cos(60)/cos(80)1/2 ~ 1.7, we can conclude that we will fall short by a factor of at least 1.7; goingto 30 latitude may not be enough for measuring meridional flow close to the pole.

Figure 31. (left:) The path of ray-path end points when a ray with the skip angle of 45, centeredaround a point at 30 latitude(60 co-latitude) is rotated. (right:) The projection factors at the end
points (dashed) as function of , the angle of the path to the local meridian, and its product which isnearly constant (solid)

67

Appendix B: Jupiter Option

Jupiter option that is described in Section 3.2 cannot reach high latitudes before the solar maximum
near the early 2020’s. However, the option gives an opportunity to go to high latitudes withincurrent technology if the kick-motor for H-IIA launch vehcle is availabe. When a longer mission
duration than the current Jupitor option is acceptable, there are solutions of spacecraft orbit with ahigher inclination. The Jupiter option whose final orbit is similar to that of the SEP option in
Sections 3.2 and 3.3 is briefly inroduced here.

B.1 Orbit Trajectory
Figure 32 shows the orbit trajectory for the Jupiter option that satisfy both the short cruise time andthe maximum orbit inclination. This is essentially a balistic orbit using conventional technology.
Minimum distance to the Jupiter is ~60 RJ, so that the radiation environment near the Jupiter is notsevere.

Figure 32. Trajectory profile of Jupiter option

B.2 Mission Profile
The sequence of events of Jupiter option is shown in Table 12. In this sequence, one Jupiterswingby and two Earth swingby operations form the final circular orbit of 1 AU distance from theSun. A similar orbit is possible for the launch opportunity every ~1.1 year but the maximum orbit
inclination is slitly diferent. In the case of 2019 launch, the maximum orbit inclination to the solar
ecliptic plane is 36.6

68

Table 12. Sequence of events of Jupiter Option
Date

Event

v

iSEQ

2018/01/12
2020/02/05
2022/12/30
2024/12/30

LaunchJupiter Swingby
Earth Swingby
Earth Swingby

8.9 km/s
7.3 km/s
17.0 km/s
17.0 km/s

7.1
31.4
35.6
39.5

Figure 33. Mission profile of Jupiter option (to be replaced into English version)

Figure 33 shows the mission profile of Jupiter option. This indicates that the first high-data rate
observations starts > 4.5 years after the launch. It takes 7 years to reach the final orbit.

B.3 Mass and Power Budget
The mass and power budget are shown in Table 13 and Table 14, respectively.

69

Table 13. Spacecraft mass budget of Jupiter Option

Table 14. Spacecraft power budget of Jupiter option
Power (W)
Item
Near Jupiter
1 AU from Sun
Mission payload
α (TBD; heater)
150
Data hadling system
21.8
21.8
Communication system
79.0
211.3
Power supply system
37.0
51.0
Attitude control system
99.0
106.7
Chemical thruster system
250.0
4.0
Total Power486.8+α70

544.8

Appendix C: References
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Cranmer, S. R. et al. 1999, ApJ, 511, 481
Dikpati, M. and Gilman, P. A. 2006, ApJ, 649, 498
Forbes, T. G. 2000, JGR, 105, 23153
Gabriel, A. H., 1971, Solar Phys., 21, 392
Giles, P. M., Duvall Jr, T. L., Scherrer, P. H. & Bogart, R. S. 1997, Nature, 390, 52
Haber, D. A. et al. 2002, ApJ, 570, 855
Harrison, R. A. et al. 2009, Solar Phys., 256, 219
Kamio, S. et al. 2009, A&A, 502, 345
Knaack, R. et al. 2001, A&A, 376, 1080
Kohl, J. et al. 1980, ApJ, 241, L117
Kohl, J. et al. 2006, Astron. Astrophys., Rev., 13, 31
Leighton, R. B. et al. 1962 ApJ, 135, 474
Lockwood, G. W., Skiff, B. A., Baliunas, S. L., and Radick, R. R. 1992, Nature, 360, 653
Lockwood, G. W., Skiff, B. A., and Radick, R. R. 1997, ApJ, 485, 789
McComas, D. J. et al. 2003, GRL, 30, 1517
McComas, D. J., Ebert, R. W., Elliott, H. A., Goldstein, B. E., Gosling, J. T., Schwadron, N. A.,
and R. M. Skoug, R. M. 2008, GRL, 35, L18103
Nakagawa, A. 2007, ApJ, 660, 1660
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Raouafi, N. E. and Solanki, K. 2006, A&A, 445, 735
Reames, D. V. 1999, Space Sci. Rev., 90, 413
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71

Appendix D: Acronyms
ACE
ACS
AR
ARD
BAT
CCD
CFRP
CH
CIR
CME
CMOS
CZ
DHU
DIM
DPCM
DSN
EAI
EDVEGA
EGA
EIS
EM
ESA
ESS
EUV
EUVI
FOV
FWHM
GCR
GONG
HAI
HCE
HK
HMI
HS
ICME
IEM
IES
IHI
IPM
IPS
IPU
ISAS
ISP
ITA

Advanced Composition Exploler (spacecraft)
Attitude Control System
Active Region
Aerospace Research and development Directorate
Battery
Charge Coupled Device
Carbon Fiber Reinforced Plastic
Coronal Hole
Co-rotation Interaction Region
Coronal Mass Ejection
Complementary Metal Oxide Semiconductor
Convection Zone
Data Handling Unit
Direct Inclining Method
Differential Pulse Code Modulation (,a type of digital communications method)
Deep Space Network
EUV Activity Imager proposed for SOLAR-C Plan-A
Electric Propulsion Delta-V Earth Gravity Assist
Earth Gravity Assist
EUV Imaging Spectrometer (on Hinode spacecraft)
Engineering Model
European Space Agency
EUV Scanning Spectrometer proposed for SOLAR-C Plan-A
Extreme UltraViolet
EUV Imager (on STEREO spececrafts)
Field Of View
Full Width at Half Maximum
Galactic Cosmic Rays
Global Oscillation Network Group (by observing the Sun with six ground-based
observatories for helioseismology)
Helioseismic Activity Imagier proposed for SOLAR-C Plan-A
HouseKeeing Electronics
HouseKeeping
Helioseismic Magnetic Imager (on Solar Dynamics Observatory)
HelioSeismology
Interplanetary Coronal Mass Ejection a type of digital communications method
Ion Engine Module
Ion Engine System
Inner Heliospheric Imager proposed for SOLAR-C Plan-A
IES Pointing Mechanism
InterPlanetary Scintillation
IES Power Unit
Institute of Space and Astronautical Science
SPecific Impulse
Ion Thruster Assembly
72

ITCU
JAXA
JGA
JPEG
JPL
LASP
LIM
MDI
MHD
MLI
MMO
MSU
NAOJ
NASA
PCU
PMOD
PMU
PREMOS
PYC
QS
RW
RZ
SAP
SDO
SEM
SEP
SOHO
SORCE
SOT
SSR
STEREO
TCIM
TIM
TR
TRL
TSI
UV
UVCS
VGA
VIRGO
XRT

Ion Thruster Control Unit
Japanese Aerospace Exploration Agency
Jupiter Gravity Assist
Joint Photographic Experts Group (, a group which created a data compression)
Jet Propulsion Laboratory
Laboratory for Atmospheric and Space Physics
Luminosity and Irradiance Monitor proposed for SOLAR-C Plan-A
Michelson Doppler Imager (on SOHO)
Magneto-HydroDynamics
Multi-Layer Insulator
Mercury Magnetospheric Orbiter
Microwave Supply Unit
National Astronomical Observatory of Japan
National Aereopace Agency
Power Control Unit
Physikalisch-Meteorologisches Observatorium Davos
Propellant Management Unit
PREcision MOnitor Sensor on PICARD spacecraft for monitring TSI
Pyrotechnic Controller
Quiet Sun
Reaction Wheel
Radiative Zone
Solar Array Panel
Soalr Dynamics Observatory (spacecraft)
Sequential EDVEGA Method
Solar Energetic Particle in science part or Solar Electric Propulsion in engineering
Solar and Heliospheric Observatory (spacecraft)
SOlar Radiation and Climate Experiment (spacecraft)
Solar Optical Telescope on Hinode spacecraft
Serieis Switching Regulator
Solar TErrestrial RElations Observatory (spacecrafts)
Telemetry Command Interface Module
Total Irradiance Monitor (on SORCE)
Transition Region
Technical Readiness Level
Total Solar Irradiance (on SORCE spacecraft)
UltraViolet
UltraViolet Coronagraph Spectrometer (on SOHO)
Venus Gravity Assist
Variability of solar IRradiance and Gravity Oscillations (on SOHO)
X-Ray Telescope (on Hinode spacecraft)73