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).

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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.
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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
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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.
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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)

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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.

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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.

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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.
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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