Chapter 1:

1.1.
1.1.1.

Science Goals and Requirements

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

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

1-11.1.2.

Elementary Structures of Coronal Magnetic Loops

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

Figure 1.1. A magnetically active region at the edge of the Sun shows "loops" that connect one
magnetic polarity to the other. A coronal image (Fe IX in 171A) from TRACE. Coronal
structures seen with 0.5 arcsec spatial resolution.
The EUV spectrometer with the spatial resolution of ~0.3 arcsec, which is about one order ofmagnitude higher than EIS resolution, (and the normal incidence X-ray telescope with thespatial resolution of 0.2 arcsec as the second option), will be able to almost resolve coronalloops at fundamental scales. That is, spatial scales at which loops behave coherently and can be
described by simple hydrodynamic models. The combination of the filling-factor analysis with
the EUV spectrometer and coronal images at 0.1 arcsec pixel from the X-ray imaging telescope
makes clear what the elementary structure of coronal loops is. For this purpose, the X-ray
telescope with normal incidence optics may be more suitable than the photon-counting
telescope.
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Figure 1.2. A polar coronal hole bright point in lines of Fe XI (red), Fe XII (green), and Fe XV(blue). The bright point looks like an unresolved flare seen by Yohkoh.

1.1.3.

Elementary Structures in Dynamical Chromosphere

The chromosphere is the region that interfaces the gas-pressure-dominated photosphere andthe magnetic-field-pressure-dominated corona, through which ionization fraction and theplasma beta change drastically. The chrompsphere forms the conduit for the upward flux ofmass and energy, while most of the non-radiative energy is dissipated there to supply the heat.
Hinode observations demonstrate that the chromosphere is extremely dynamic: small scale
structures called as straws, jets, or fibrils cover the entire surface of the Sun, high speed
chromospheric ‘jets’ (type-II spicules, de Pontieu etal 2007) have a width of ~0.3arcsec andchange with a time scale of 10 second or less. It is obvious that the chromosphere cannot be
considered as an atmospheric ‘layer’, but it is a wholly dynamic 3D entity that never settles tohydrostatic equilibrium. Even in the most quiet area of the Sun, i.e., in the very quiet
photospheric “internetwork”, the chromosphere appears to be dominated by largely acoustic
shocks, which are initiated in the photosphere by the solar convection and the global p-mode
oscillations. Strong “network” fields exist in the photosphere at the boundaries of theinternetwork. These fields spread rapidly with height in the highly stratified chromosphere,
and at some height they form a magnetic canopy over the internetwork regions. Recent
observations from Hinode and other instruments have suggested that the network fields,interacting with convection, energize and structure the quiet chromosphere. Dynamic
phenomena in the chromosphere are more dominating in active regions, where apparent
reconnection events occur frequently, producing outward propagating jets that may be seen atcoronal temperatures (e.g., Shibata et al. 2007, Figure 1.3).
High-speed jets in quiet Sun (type-II spicules) are, however, only visible near the solar limb
with Hinode imaging observations, because of their insufficient optical thickness on the disk
under the Hinode’s Ca II H band-pass filter. In addition, Hinode provides only morphologicalinformation of such chromospheric features. With the lack of the Doppler information, thenature of the type-II spicules in quiet regions or those in sunspot penumbrae (Katsukawa et al.
2007) are still under debate. One of important question is whether they are real plasmaejections with the apparent speed along the magnetic fields or not. Without answering thequestion, we cannot know their physical origin. To understand the generation mechanism andconsequence of this fundamental ingredient of the chromospheres, we definitely need velocity
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information with better contrast and higher spatial and temporal resolutions. The proposed
Solar Ultra-Violet Visible and IR telescope (SUVIT) has a large aperture (1.5 m diameter is
now under study), which can achieve a spatial resolution better than 0.06arcsec in Mg II
280nm line. This UV line provides unprecedented images of fine-scale chromosphericstructures with much higher contrast. The supreme spatial resolution allows us to explore
plasma motions and magnetic-field behaviors inside the chromospheric structures with
~0.3arcsec width. Observations of full spectral line profiles allow us to study the real plasma
flow along the magnetic fields with its acceleration and density distribution, and also toclearly identify the connectivity to their photospheric (or possibly chromospheric) root points.
The high sensitivity, high-resolution Doppler images (Dopplergram) of chromospheres will
for the first time detect the elementary plasma motions, which would be associated with wave
shocks and magnetic reconnections, and thus provide a concrete basis of energetic in thechromosphere.

Figure 1.3. Dynamical chromosphere observed with Hinode. (Left) Dynamical chromosphericfeatures around a large sunspot (Right) Chromospheric jets.
Ubiquitous chromospheric jets associated with cusp-shaped structures (Figure 1.3) show ahint of intermittency of plasma ejection, which infers the presence of more elementary
reconnection processes. SUVIT aims to identify any non-stationary motions or intensity
fluctuation in and around the reconnection site to understand the basic physics that controls
the rate of fast reconnection taking place in the chromospheres (See also Section 1.3). One
particular interest regarding the magnetic reconnection is the injection of magnetic helicity
into the magnetic fields during the reconnection of obliquely directed two flux tubes.
Spectroscopic information in a higher spatial resolution will detect the twist or twisting
motions in a flux tube, which may be crucial importance for dynamics of the ejected plasma
or even for global evolution of the corona. Moreover, SUVIT provides the direct
measurements of the magnetic fields in particular chromospheric features like thereconnection cusp. Even with less spatial and temporal resolution in polarimetric observations,such direct information of magnetic fields at the reconnection site is of a fundamental
importance to establish the reconnection models of the chromospheric dynamics. Enigmatic
phenomena, such as the Ellerman bombs around sunspots (Matsumoto et al. 2008) andintermittent and recurrent occurrence of chromospheric plasma ejections from sunspot lightbridges (Shimizu et al. 2009), are also important targets with a great interest relating to themagnetic field reconnection.
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1.1.4.

Elementary Structures in Photosphere

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

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

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

Figure 1.6. The sensitivity to the magnetic flux and energy elements (εdx2) and the field of view(FOV) of SUVIT and other observations. The SUVIT sensitivity is compared with Hinode SOT
and ground-based observations. The unit in the vertical axis is arbitrary and the horizontal axis
in arcsec. The figure also shows the sensitivity range and FOV of the observations required for
4 specific science topics.
A similar argument can be applied to the Doppler measurements too; i.e., SUVIT has a 10
times higher sensitivity than SOT to the kinetic energy of gas motions. Tiny flows and shocks
possibly associated with photospheric magnetic reconnections, but not resolved so far, are one
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of the important targets of SUVIT observations. Observing the dynamic interactions of themagnetic fields with ubiquitous shocks generated by granular convection in the photospheremay provide a direct clue on the origin of spicules or transient phenomena in chromosphere.
These fundamental plasma processes postulated to be taking place in the high β photosphericplasma are only accessible with high accuracy spectro-polarimetric measurement at 0.1"
resolution under a very stable imaging condition, which can be only realized from the space.

1.2.

Understanding Energy Transport through Small-scale Structures

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

Connecting the Chromosphere to the Corona

Hinode imaging observations revealed that substantial transverse displacements exist in
many of chromospheric “Type-II” spicules (De Pontieu et al. 2007). Moreover, withcoordination with measurements of EUV spectral lines oriented from the transition region (TR),a blue-red asymmetry was derived in TR spectral lines, suggesting that “Type-II” spicules
transport mass towards the corona (De Pontieu et al. 2009, McIntosh & De Pontieu 2009).
Furthermore, the spectral lines oriented from the corona at temperatures of 2MK show excess
non-thermal line broadening at the footpoints of coronal loops, which can be seen only when
the coronal loops are observed from the top (i.e., disk center view) (Hara et al. 2008). Based onthese Hinode observations (Figure 1.7), quantitatively evaluating of mass and energy transport
between the chromosphere, TR, and the corona at the base of coronal loops is essential to reachthe conclusion for the coronal heating puzzle.
TR (transition region) is the interface connecting between 20,000 K (upper chromosphere)
and 1 MK (corona). SOHO, Hinode and previous space missions have shown the TR to be ahighly dynamic and structured region. SUMER high-resolution spectro-heliograms reveal that,on the quiet Sun, the lower-mid TR (T < 700,000 K) is formed by many loop-like structures 10″
to 20″ long located along and across network boundaries. Similar but fainter structures are also
seen in cell centers. From these data it seems that the majority of the TR emission, particularlyon the quiet Sun, may not represent a continuous transition between the chromosphere and thecorona. EIS raster scans also show the transition from the small-scale TR structures tolarge-scale coronal loops and funnels to occur around 0.7 to 0.8 MK (Matsuzaki et al. 2007).
Of course, transition region “interface” regions must exist at the footpoints of loops or
funnels reaching coronal temperatures but they appear to be responsible for only a small
fraction of the TR emission. This would naturally explain the fact that the TR emits much morethan it would do if it were a very thin (few tens of km) thermal layer, heated by thermal
conduction from the hot corona (as in classical, static, 1-D models). On the other side, SOHO
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observations have also shown the TR to be very dynamic on time scales of the order of 10 s to
100 s.
Quiet Sun TR loops also do not seem to have obvious magnetic counterparts at photosphericlevels, indicating that the loop footpoints are connected to multiple magnetic elements in thephotosphere via a complex system of field lines (Sanchez-Almeida et al. 2007). A similar
phenomenon is also observed in moss areas, regions of enhanced EUV emission in upper TR
and low coronal lines at the footpoints of 3 MK to 5 MK coronal loops in active regions. In thiscase their emission is mainly due to thermal conduction from the hot loops. Each footpoint ofcoronal loops is rooted into a uni-polar magnetic plage region, where no mixed polarity
magnetic flux elements are observed at the photospheric level, suggesting a complex magnetic
coupling in the upper atmosphere (Brooks et al. 2009).
Therefore, measuring magnetic fields in the higher atmosphere, i.e., chromosphere, would be
crucially important for unveiling the nature of this complex coupling between loop-likestructures in TR and the photosphere. At the same time, the EUV spectrometer probes all
temperature ranges from the upper chromosphere up to the corona with high spatial, spectral,
and temporal resolution. The TR strongly emits in the UV/EUV (160 to 2,000 Å) and lines fromincreasingly higher ionization stages generally fall at shorter wavelengths. The EUV
spectrometer should cover a very wide spectral range at unprecedented spatial and temporal
resolution, allowing, for the first time, the full characterization of this region of the solar
atmosphere: from the topology and morphology of the magnetic structures to thethermodynamic state of the plasma (its motions, density, temperature and abundances) that fills
these structures.

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

1.2.2.

Magnetohydrodynamic (MHD) Waves in the Solar Atmosphere

Recent observations, especially Hinode’s imaging observations, have clearly revealed thepresence of apparent transverse oscillations, which are affirmative evidence of the Alfven
waves, in the chromosphere and corona. Wave-like oscillations are observed in a wide variety
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of solar structures, such as prominence in the corona (Okamoto et al. 2007, Figure 1.8), spicules(de Pontieu et al. 2007, Figure 1.9), and Ca II H dynamical jet (Nishizuka et al. 2008) with
Hinode SOT.

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

Figure 1.9. Example of the transverse displacement of a spicule (from de Pontieu et al. 2007)
However, we did not identify whether the observed waves are really propagating Alfven
waves, or just standing waves in most of observations, except for ubiquitous upward waves in
the corona detected with a coronagraph measurement of line-of-sight velocity fluctuations(Tomczyk et al. 2007). Thus, in order to evaluate the propagating MHD waves and their modes,
observations require not only velocity and photometric measurements (Kitagawa et al 2010)
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but also the measurements of the magnetic field fluctuation, as examined with photosphericmagnetic-field data by Fujimura and Tsuneta (2009).
1.2.2.1. Relevance of Observed Waves for Chromospheric Energy Budget
The chromosphere’s temperature is approximately 10,000 K with high density. To maintainthe temperature structure of the chromosphere, it is required to input heat in the chromosphereabout one order of magnitude higher than that required for heating the corona.
Recent studies show that there is insufficient acoustic wave power propagating at
higher-frequencies (>5mHz) from the solar interior to heat the chromosphere (Fossum &Carlsson 2005; Carlsson et al. 2007). However, it is shown that inclined magnetic field lines atthe boundaries of large-scale convective cells provide portals through which low-frequency(<5mHz) magnetoacoustic waves can propagate in the chromosphere (Suematsu 1990), andthese waves provide a significant source of the energy necessary for balancing the radiative
losses of the ambient chromosphere (Jefferies et al. 2006). When magnetoacoustic waves
propagate upward in strongly stratified atmosphere such as the chromosphere, the waves will
steep into shocks to heat the chromosphere (Carlsson & Stein 1997). Carlsson, Hansteen &Gudiksen (2010) indicates that as low-frequency magneto-acoustic shock waves interact withsmall magnetic elements, joule heating may be effective at these elements’ edges.
It is important to observationally determine how much energy is sustained in observed waves,with identifying the mode of the waves. The energy should be measured as a function ofmagnetic field inclination and atmospheric height. To determine the power of magnetoacoustic
waves at the lower atmosphere, we need the measurements in cadences of better than 10-20
seconds with line-of-sight velocity amplitudes of ~0.1 km/s for the photospheric lines and with
~1 km/s for chromospheric lines (e.g., Vecchio et al. 2007). High spatial resolution, 0.1-0.2
arcsec, is needed to resolve individual channels for wave propagation (Bogdan et al. 2003).
Moreover, determining the importance of shocks is vital for balancing the radiative losses in the
more magnetized parts of the upper chromosphere. To do this, velocity measurements with UVlines in formation temperatures between 104 and 105 K need to be coordinated with thephotospheric and chromospheric measurements with visible lines. For measurements with UVlines, we need cadences of better than 20-30 seconds, which provide 6-15 measurements in one
period of 3-5 minutes oscillations, and Doppler velocity measurements of better than 10 km/s
(equivalent to sound speed in the upper chromosphere) with line width measurements of better
than 5 km/s. It is also important to have the capability to resolve the elemental structures where
the wave power is channeled. These are typically on the order of 0.3-0.4” in size, inferred fromthe observed width of spicules. The wide temperature coverage will allow us to track theevolution of the wave power as it propagates through the different layers of the solar
atmosphere. We should note that the NASA IRIS mission, which will be launched in late 2012,
will provide new information on the chromosphere, especially on the line-of-sight velocity
amplitude measured with chromospheric lines, but it does not have adequate transition region
and coronal coverage, nor will it measure chromospheric magnetic fields, both of them areessential to evaluate the natures of magnetoacoustic waves in the solar atmosphere.
1.2.2.2. Time Series Spectra at Different Atmospheric Layers
The problems encountered in understanding the coupling between the chromosphere and thecorona are illustrated in Figure 2.2-4. The chromosphere has relatively large patches (10′′ to20′′) of many thin, bright spicules (width~0.3”), and the coronal structures are narrow, isolated
loops that fade at their footpoints. There is little correspondence between the two. A
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comparable transition region image that can only be obtained by a rastering spectrometer would
take almost two hours to complete with the current available spectrometer such as SUMER.
Although the large-scale structures do not normally change on this timescale, the small-scale
dynamics are much faster, and thus high-cadence spectroscopic measurements in transition
region lines will make a giant step forward. Moreover, matching spatial resolutions both below(chromosphere and photosphere) and above (transition region and corona) is crucially
important to trace the energy flow over the entire atmosphere.
As illustrated in Figure 1.10, spectral lines show different characteristics among
chromosphere (Cauzzi et al. 2007), transition region, and corona. The chromospheric spectra
have characteristic 3 to 5 min periods, the transition region brightenings occur in bursts of up to
30 min (Innes et al. 1997), while coronal variations are more diffuse. Such different
characteristics would reflect the energy transport and release at each layer, but no one knows
details, because this kind of time series spectra has never been recorded simultaneously withdifferent spectral lines covering the solar atmosphere with 0.3” spatial resolution.

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

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

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

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

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

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

1-14

1.3.
1.3.1.

Magnetic Reconnection
The Solar Atmosphere as a Laboratory for Magnetic Reconnection

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

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

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

1.3.2.

Magnetic reconnection occurring in a wide variety of plasma conditions

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

Gas
NeutralCollisional

Chromosphere

Inner Corona

Magnetic field
Weakly ionized
Collisional

Magnetic field
Fully ionizedSemi-collisionless

1-16

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

Magnetic reconnection in weakly ionized plasma (solar chromosphere)

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

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

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

Magnetic reconnection in collisonless plasma (corona and transition region)

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

amongkinetic,

thermal,

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

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

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

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

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

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

Particle Acceleration coupled with Reconnection Dynamics

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

1-21

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

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

1-221.3.6.

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

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

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

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

Flares and CMEs

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

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

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

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

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

1-26

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

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

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1.4.2.

Solar Winds

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

Figure 1.21. Some possible sources of the fast solar wind (plume, interplume, jet) and slow
solar wind (coronal hole-steamer boundary, edge of active regions, reconnection between open
field lines and close loops, reconnection at streamer stalk).
Being the most prominent feature within CHs, plumes have been and are considered as astrong candidate for the source regions of the fast wind, but theoretical models give opposite
indications; Del Zanna et al. (1997) and Casalbuoni et al. (1999) show that the outflow speed inplumes can be larger or smaller than in the ambient depending on the temperature and Alfvén
wave flux assumed in the two regions. VUV lines formed around 1 MK allow the study of theplumes roots on the disk where their intensity is enhanced with respect to the background CH
emission. In addition to measuring Doppler shift along line of sight, Doppler dimming features
seen in, for example, O VI 103.2/O VI 1037A lines allows to determine the radial component of
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the plasma outflow velocity from the analysis of off-limb spectra, although some assumption
on the geometry of coronal structures is necessary. From SUMER and UVCS combined data,
Teriaca et al. (2003) derived the outflow vs. altitude profile in interplumes. However, the result
for the plumes is controversial with very different results being present in the literature(Giordano et al. 2000; Gabriel et al. 2003, 2005; Raouafi et al. 2007). Jets are often associatedto plums (Raouafi et al. 2008), which may cause different results, because the integration of
signals recorded by the 36 to 48 hours’ SUMER observations is required to derive reliable
results. High throughput performance improved by one order of magnitude from SUMER willallow us to study time variability of outflows at timescale of a few hours. From EUV
spectroscopy, plumes are known to be denser and cooler than the surrounding interplumeregions (e.g., Wilhelm 2006). Spectral lines are observed to be broader in interplumes (e.g.,
Wilhelm, 2006; Banerjee et al. 2000) hinting at preferential energy deposition in thesestructures. Signatures of slow magnetoacoustic waves have also been found in both plumes(DeForest & Gurman 1998, Ofman et al. 1997; 2000) and interplume regions (Banerjee et al.
2000; 2009), and it is important to understand observational link between waves and velocity
structures (see also section 1.2). Moreover, the spectroscopic analysis of chemical composition
could provide a tool to establish whether plumes or interplume regions contribute significantly
to the fast wind by comparing the composition of the different structures on the Sun with in situ
data.

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

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

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

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

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