Concept Study Report
概念検討書

SOLAR-C WG
Date 29, Jan, 2018, version 1
1. Science goals of the mission [1-1-1]
`YUbgFAs a fundamental step towards answering how the plasma universe is created and evolves, and how
the Sun influences the Earth and other planets in our solar system, the proposed mission is designed
to comprehensively understand how mass and energy are transferred throughout the solar
atmosphere. Understanding the solar atmosphere, which connects to the heliosphere via radiation,
the solar wind and coronal mass ejections, and energetic particles is pivotal for establishing theconditions for life and habitability in the solar system.
2.

Scientific objectives of the mission [1-2-1]
`YUbgF( &Solar-C_EUVST (EUV High-Throughput Spectroscopic Telescope) is a mission designed toprovide a conclusive answer to the most fundamental question in solar physics: how does theinterplay of magnetic fields and plasma drive solar activity (Figure 2.1)? The most significant
examples of this interplay are atmospheric heating and explosive energy release, such as flares andcoronal mass ejections (CMEs). Thus, the two primary science objectives for Solar-C_EUVST are:I. Understand how fundamental processes lead to the formation of the solar atmosphere andthe solar wind.
II. Understand how the solar atmosphere becomes unstable, releasing the energy that drives
solar flares and eruptions.
To achieve these science objectives Solar-C_EUVST will
A. Seamlessly observe all the temperature regimes of the solar atmosphere from thechromosphere to the corona at the same time,
B. Resolve elemental structures of the solar atmosphere with high spatial resolution andcadence to track their evolution, and,C. Obtain spectroscopic information on the dynamics of elementary processes taking place inthe solar atmosphere.
Recent missions such as Hinode (Solar-B) have clearly demonstrated that different layers of thesolar atmosphere are highly coupled with each other via magnetic fields (see Figure 2.2). This
interplay happens at spatial and temporal scales that have not been resolved over the full range oftemperature extremes 0.01 MK in the chromosphere to 1 MK in the corona to 10 MK in flares. To
understand the physical interplay between the plasma and magnetic field, we must determine how
mass and energy are transported, stored, and converted into other forms of energy in the solar
atmosphere over vastly different physical domains.

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Several recent missions have given a glimpse into
what can be observed in the solar atmosphere athigh spatial and temporal resolution. For example,
during its 300 s rocket flight in 2012, the High
Resolution Coronal Imager (Hi-C; Kobayashi et al.
2014) observed the solar corona at a resolution of
0.3-0.4" and a cadence of 5.5 s and was able todetect heating events that were largely invisible tolower resolution instruments (Testa et al. 2013).
Furthermore, Hi-C also observed very narrow,braided field lines that appeared to be reconnecting
and producing high-temperature plasma (Cirtain etal. 2013).
Finally, to truly understand the interplay ofmagnetic fields and plasma, simultaneous Figure 2.1: Solar coronal structures and activities inobservations of the photosphere are essential. soft-Xray.
Several previous missions, such as SoHO, Hinode,and SDO (see section 3.3), have carried instruments to observe both the solar photosphere and theupper atmosphere. The latest generation of ground-based optical telescopes, however, represent aquantum leap over previous space-based instrumentation. Foremost among these observatories is
the Daniel K. Inouye Solar Telescope (DKIST), which will have a diameter of 4 m and achieve anangular resolution as high as 0.03" or about 20 km on the solar surface. DKIST is nearing
completion and is scheduled to become fully operational before the end of 2020. In 2027, the 4m
European Solar Telescope (EST) is also planned to start observation, doubling the opportunity for
coordinated observations of the magnetic field.	
  To achieve the final link between the plasma
dynamics and the magnetic field, Solar-C_EUVST observations will be coordinated with DKIST
and other ground-based observatories located throughout the world.

Figure 2.2: The layers of the solar atmosphere are coupled to each other by magnetic fields.

2/100For the main scientific objectives we define sub-objectives (I-1~I-4, II-1~II-2) to be addressed with
the Solar-C_EUVST mission. To tackle the sub-objectives, we also define the observing tasks atright side (labeled such as I-1-1) of Table 2.1, which describe what we want to observe with themission. Details of the observing tasks are described in sections 3.2 and 5.1.
Table 2.1: Science objectives of the Solar-C_EUVST mission
I. Understand How Fundamental Processes Lead to the Formation of the Solar Atmosphere and theSolar Wind
Historically, two primary mechanisms have been proposed to explain how the chromosphere and corona areheated, namely small-scale magnetic reconnection and wave dissipation. Recently the role of small
chromospheric jets called spicules has also been studied extensively. These mechanisms are also directly
linked to the origin and acceleration of the solar wind. By investigating energy transfer and release on small
spatial and temporal scales, Solar-C_EUVST will quantify the relative contributions of these mechanisms to
the formation of the solar atmosphere and the solar wind.
I-1
Quantify the Contribution of Nanoflares to Coronal Heating
Nanoflares,
small-scale
magnetic I-1-1
Measure the energy of small-scale heatingreconnection events, are a possible
events in the transition region and the corona inmechanism for heating the corona. Solarthe energy range of ~1024 - 1027 erg.
C_EUVST will evaluate this hypothesis I-1-2
Observe intermittent processes that generateby observing high-temperature plasmas
plasmas above 5 MK with high speed plasma
and their dynamical behavior throughmotions.
resolving elemental structures in the I-1-3
Observe sub-arcsec braiding structures with highcorona.
temporal and spatial resolutions.
I-1-4
Identify the driver of nanoflares by comparing
spectroscopic diagnostics with simultaneous
observations of the photosphere and low
chromosphere.
I-2
Quantify the Contribution of Wave Dissipation to Coronal Heating
The wave heating hypothesis suggests I-2-1
Detect Alfvén waves by measuring thethat waves propagate upwards from the
propagation of fluctuations through different
solar surface and are dissipated, leading
layers of the atmosphere.
to the heating of the solar atmosphere. I-2-2
Observe the thermalization process by
Solar-C_EUVST will quantify thismeasuring how transition region and coronal
scenario by measuring the characteristics
plasmas respond to the propagating waves.
of the waves at different heights and I-2-3
Identify the source of upwardly propagating
observing the thermalization process.
waves by comparing spectroscopic diagnostics
with simultaneous observations of thephotosphere and low chromosphere.
I-3
Understand the Formation Mechanism of Spicules and Quantify Their Contribution to Coronal
Heating
A spicule is a dynamic jet launched I-3-1
Observe the thermal evolution of spicules (width
upwards from the lower atmosphere and~0.4") from chromospheric to transition region
is a fundamental ingredient of the solar
and coronal temperatures. Quantify the mass
chromosphere and the transition region.
flux that spicules supply to higher altitudes.
Solar-C_EUVST will clarify how I-3-2
Identify the driving mechanisms of spicules by
spicules are created and quantify theircomparing spectroscopic diagnostics withmass and energy contribution to coronalsimultaneous observations of the photosphereheating.
and low chromosphere.

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I-4

Understand the Source Regions and the Acceleration Mechanism of the Solar Wind
The solar wind is the plasma flowing I-4-1
Observe the velocity, temperature and density
along open field lines into thestructures at the source regions of solar wind andheliosphere. This plasma is thought to beclarify their relationship to the magnetic field
accelerated by Alfvén waves as it flows
structures.
away from the solar surface. The wind I-4-2
Detect signatures of coronal Alfvén waves inoriginates in faint regions of the corona,
plume and inter-plume regions and measure
making it difficult to measure the plasma
their energy fluxes with height.
properties. Solar-C_EUVST will provide
sensitive measurements of the velocity,temperature, density, and abundance inthese regions, revealing the formation
and acceleration mechanisms of the solar
wind.
II. Understand How the Solar Atmosphere Becomes Unstable, Releasing the Energy that Drives Solar
Flares and Eruptions
Photospheric motions lead to the accumulation of free magnetic energy in the corona. This system
eventually becomes unstable, releasing the energy through magnetic reconnection. This process of energy
conversion heats the plasma to high temperatures and drives coronal mass ejections (CMEs). By measuring
the properties of multi-temperature flaring plasma, Solar-C_EUVST will investigate why the reconnection isfast despite the high magnetic Reynolds number. It will also monitor the temporal evolution of solar active
regions and identify the triggering mechanism for the flare and eruption.
II-1 Understand the Fast Magnetic Reconnection Process
Magnetic reconnection is one of the II-1-1 Probe plasma conditions and structures inside
fundamental processes for converting
the reconnection region and clarify the role ofmagnetic energy into the thermal andshocks and magnetic islands in fast
kinetic energy of the plasma. This
reconnection.
process occurs much faster than is II-1-2 Probe the conversion of energy by observing thepredicted by classical theory. Solarchromospheric response to magnetic
C_EUVST will observe the dynamics ofreconnection at very high cadence.
magnetic structures to understand the II-1-3 Characterize the physical properties andmechanisms that lead to fast magnetic
dynamics of magnetic reconnection occurring inreconnection in partially or fully ionized
the chromosphere and transition region, whereplasmas.
the plasma is different from the fully ionizedplasma of the corona.
II-2 Identify the Signatures of Global Energy Buildup and the Local Triggering of the Flare andEruption
Understanding the accumulation and II-2-1 Monitor long-term, large-scale evolution of
release of free magnetic energy in theactive regions and identify the spectroscopic
corona is a fundamental problem. Solarsignatures of energy buildup.
C_EUVST will perform long-term II-2-2 Characterize the dynamics of small-scalemonitoring of active regions to identify
magnetic structures that trigger the eruption of
the signatures of energy buildup andflaresandidentifytheMHD
high-resolution
observations
to(magnetohydrodynamic) instability modes by
understand the triggers of energy release.
comparing photospheric and low-chromosphericobservations against numerical modeling.

3.
3.1.

Rationale for the scientific goals and objectives
`YUbgF( -B&F
Scientific background, status of the relevant scientific area, etc. ( %-Fk')F.j14F')"!j*There are two crucial facets of solar research that underline the scientific priority of SolarC_EUVST. First, many of the phenomena that take place on the Sun also occur elsewhere in theuniverse. Therefore, the insights gained from Solar-C_EUVST will advance the knowledge of
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fundamental processes in astrophysical plasmas. Second, the Sun is the star that governs our spaceenvironment and, for a technologically advanced global society, this space environment is directly
connected to our daily lives. Increasing our knowledge of solar variability is critical for humanity
due to our ever-increasing reliance on space-based technologies for daily life. In this regard, there
are several imperative scientific investigations, including understanding large-scale solar eruptions
and the origin of the solar wind and its dynamics that directly affect our global environment anddaily lives.
To reinforce the international support for this proposal, we would like to cite a study report on the
future of solar physics for the decade of 2020 that was published by the Next Generation Solar
Physics Mission (NGSPM)’s Science Objectives Team (SOT), an advisory team consisting of 14
scientists, chartered by the United States National Aeronautics and Space Administration (NASA),
the Japan Aerospace Exploration Agency (JAXA), and the European Space Agency (ESA). The
report (see Supplement A) lists scientific objectives in all the research areas related to solar physics
and provides priorities for next generation solar physics missions to be realized in the mid 2020s.
The report is based on extensive reviews of the broad interests of the heliophysics researchcommunity, following a public call for white papers in the fall of 2016 which resulted in 34
submissions describing science objectives, instruments and mission concepts, covering a wide
variety of topics. Three top-level science objectives are identified: I) Formation mechanisms of the
hot and dynamic outer solar atmosphere, II) Mechanisms of large-scale solar eruptions andfoundations for predictions, and III) Mechanisms driving the solar cycle and irradiance variation.
The Solar-C_EUVST mission aims to mainly cover topics I) and II) following a unique approachthat is described in detail below.
I) Formation mechanisms of the hot and dynamic outer solar atmosphere
Hinode has revealed the dynamic behavior of small-scale magnetic field structures (~0.4") in theform of magnetic flux tubes at the photosphere. It is clear that the turbulent motions of thephotosphere are responsible for the heating of the solar upper atmosphere. The mechanisms
responsible for transporting energy up from the photosphere and dissipating it, however, are still not
fully understood. Three plausible mechanisms, nanoflares, waves, and spicule heating, have beenstudied as possible solutions to the coronal heating problem. Furthermore, solar wind acceleration is
also often discussed in the same framework.
Nanoflares: The photospheric motions lead to the twisting and braiding of the magnetic field(Figure 3.1-A) which permeates the solar atmosphere. Parker (1972) was among the first torecognize that the dissipation of this topological complexity through magnetic reconnection could
heat the solar corona to high temperatures. Most previous spectroscopic instruments, however, havebeen unable to routinely follow small-scale heating events at the “nanoflare” size (1024-1027 erg)
envisioned in the Parker model (Winebarger et al. 2012, Ishikawa et al. 2014). Thus the detailed
plasma properties of small-scale coronal heating events are largely unexplored.
Waves: The magnetized solar atmosphere is filled with a rich assortment of waves that are an
important source of energy. Some of the oldest theories of coronal heating involve the dissipation ofwaves that are generated at low heights and propagate upward into the chromosphere and corona(e.g., Ionson 1978, Heyvaerts & Priest 1983). Imaging observations by SoHO/EIT identified thatthe coronal waves are the slow-mode MHD wave, but the total energy brought by this mode was notenough to heat the corona. Coronal Multi-Channel Polarimeter (CoMP, ground-based coronagraph)detected Alfvén waves in line-of-sight velocity images of the solar corona (Figure 3.1-B, Tomczyk
et al. 2007). Hinode has detected Alfvén waves in the chromosphere (Okamoto et al. 2007). Models
based on this concept, however, have a difficult time reproducing the observed temperatures in thecorona unless the wave fluxes are much larger than are generally observed (e.g., van Ballegooijen et5/100

al. 2017). There have been no observations with seamless temperature coverage through thechromosphere to the corona as well as high temporal resolution.
Spicules: Previous instruments such as Hinode/SOT have clearly revealed that the chromosphere isdominated by a multitude of thin (average of 300 km or 0.4"), dynamical, jet-like extrusions called
spicules (Figure 3.1-C, De Pontieu et al. 2007). Combined Hinode/SOT and IRIS observations havegreatly advanced our understanding of spicule evolution over a wider range of temperatures. It
appears that a subset of chromospheric spicules can be heated to coronal temperatures as they riseand supply hot mass to the corona or that a considerable amount of energy could be channeled into
the corona by small-scale eruptions driven by photospheric vortices or cyclones (e.g., WedemeyerBöhm et al. 2012, Iijima & Yokoyama 2017). Existing spectroscopic observations at higher
temperatures, however, do not have the spatial resolution or cadence to follow spicule evolution
further and determine how much they contribute to the mass and energy budget of the corona. Evenat low resolution, the measurements of asymmetric line profiles suggest that the contribution of
spicules is important (De Pontieu et al. 2009).

Figure 3.1: Cartoon of the formation mechanisms of the hot and dynamic outer solar atmosphere (bottom). A)
nanoflare with braiding structure observed by the Hi-C rocket, B) propagating Alfvén wave observed by CoMP, C)
spicules observed by Hinode/SOT, D) the origin of the slow solar wind at the edge of an active region observed byHinode/EIS.

Solar wind: A significant fraction of the magnetic field on the Sun is open and extends into theheliosphere. Thus, some of the same processes that lead to the formation of high temperature loops
in the closed corona heat and accelerate the solar wind. In situ observations show that the solar wind
has, at least approximately, a bi-modal nature with fast (~800 km/s) and slow (<400 km/s)
components (e.g., McComas et al. 1998). The fast wind is known to originate in the polar coronal
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holes. The origin of the slow wind, however, has not been fully established. Equatorial coronalholes, active region (AR) outflows (Figure 3.1-D, Harra et al. 2008), and streamers have all been
proposed as sites for the slow solar wind. Furthermore, the mechanism for accelerating both fastand slow solar winds have not been identified. The height at which the solar wind acceleration
begins is still controversial; some observations suggest lower heights (less than 1 solar radii, Hahn
& Savin 2013) while theory and other observations suggest higher heights (larger than 2 solar radii,
Gabriel et al. 2003, 2005). Because the solar wind acceleration region is very faint, newspectroscopic observations with high throughput are needed to detect signatures of Alfvén waves
and measure energy fluxes with height.
II) Mechanisms of large-scale solar eruptions and foundations for predictions
It is now widely accepted that solar flares are caused by magnetic reconnection, the physicalprocess that rearranges the magnetic configuration and converts magnetic energy into kinetic andthermal energies with non-thermal particle acceleration. While previous space missions have
improved our understanding of flare reconnection, there still remain several key mechanisms thatneed to be resolved, especially those related to fast magnetic reconnection. As the strongest flares,
and ones that affect Earth, often occur in the ARs, we focus on the large-scale evolution of ARs andthe triggering process of flares. These studies are key for formulating flare prediction algorithms.
Solar flares: Solar flares represent the high-energy tail of the distribution of magnetic reconnection
events and are therefore ideal for studying how magnetic energy is dissipated in the corona via
reconnection. Although the Yohkoh, SoHO, TRACE, RHESSI, Hinode, SDO, and IRIS missions
have led to a better understanding of solar flares, features such as inflows, high-velocity jets andshocks, which are thought to be central to the reconnection scenario, have not yet been observed
systematically (Figure 3.2-A, Daughton et al. 2011). It has been proposed that these features must
be too faint to observe as intensity features with current instrumentation (Figure 3.2-C). However,they should be visible as Doppler shifted features in spectra, while rapid heating at the location
where shocks are observed should result in non-equilibrium ionization and different ion andelectron temperatures. There have been some observations of these phenomena (e.g., Hara et al.
2011; Imada et al. 2013), but previous spectroscopic instruments have lacked the throughput needed
to quickly scan large areas and routinely observe flares.
Fast magnetic reconnection: Standard magnetic reconnection theory (e.g., the Sweet-Parker
reconnection model) predicts reconnection rate that is too slow to explain observed reconnection
events in the solar corona. Several theoretical studies recently suggested that small magnetic islandsin the reconnecting current sheet are important. The development of magnetic islands or plasmoidsin the current sheets (e.g. Loureiro et al. 2007, see Figure 3.2-B) can lead to fast reconnection (e.g.,
Shibayama et al. 2015). To clarify the fast magnetic reconnection process, observing thereconnection region is essential. The temperature evolution of the reconnection region can beinferred by measuring the intensities of lines of highly ionized iron, e.g. from Fe XVIII through Fe
XXIV (6 ~ 15 MK), while high velocity flows are expected in or close to the reconnection region. It
is therefore crucial to determine the current sheet geometry and test the standard models. However,this can only be done using high cadence observations with high spectral resolution of the fullgroup of hot Fe lines (Fe XII to XXIV: 1.5~15 MK) within and above the forming X-ray loops,
which is difficult with existing spectrometers as they do not possess enough throughput.
AR evolution and flare triggers: As illustrated in Figure 3.2-D, solar flares occur as a result of alarge-scale AR development, namely, the storage of free magnetic energy in the corona. In responseto the photospheric evolution, the AR corona is known to show spectroscopic signatures such as
upflows with non-thermal line broadening in the pre-flare phase (e.g. Hinode/EIS observations by
Harra et al. 2009, Imada et al. 2014). However, details of this relationship between coronalsignatures and AR evolution are not clear due to the lack of long-term monitoring with high-enough
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spatial resolution. Because flares may be triggered by small-scale magnetic structures (e.g. Bambaet al. 2013), the triggering regions should also be inspected with high-resolution spectroscopy. The
obtained data will be compared with numerical simulations to identify their MHD instability modes,
which could lead to the ability to predict flares in the near future.

Figure 3.2: Standard flare model (center: Tsuneta 1997) with observations and numerical simulations in each feature.
(A) A numerical simulation of the reconnection region (Daughton et al. 2011). (B) Erupting plasmoid andreconnecting current sheet observed by SDO/AIA (Liu et al. 2013). (C) Downward flowing loop signatures above aflare loop (Imada et al. 2013). (D) Flare eruption in an AR with flare-triggering fields (Bamba et al. 2013).

3.2.

Goals and objectives: new scientific steps the proposed concept aims to achieve in thescience area
提案するミッションコンセプトは当該分野になにをもたらそうとするのか

In order to advance our understanding of the mysterious Sun, especially of the origin of the hotsolar atmosphere and the occurrence of the solar flares, the Solar-C_EUVST mission concept
tackles the scientific objectives in section 2 by taking the following unique approaches:A.

To seamlessly observe all the temperature regimes of the atmosphere from thechromosphere to the corona simultaneously,B.

To resolve elemental structures of the solar atmosphere and track their changes with
sufficient cadence, and,C.

To obtain spectroscopic information on dynamics of elementary processes taking place inthe solar atmosphere.

This mission concept provides a completely new set of spectroscopic tools to examine the solar
atmosphere. The goals and objectives of this mission, together with a set of specific observing
tasks, are described in the rest of this subsection. These shall be the significant steps toward thecomplete understanding of thermodynamics working in the Sun.

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I-1: Nanoflares, tiny magnetic reconnection events, are a possible mechanism for heating the corona
if the total amount of energy released by nanoflares is significant. This is the reason why thedefinitive measurement of the occurrence rate for small-energy heating events is essential. SolarC_EUVST will evaluate the nanoflare hypothesis by extending the observable energy range down
to ~1024 erg, i.e., about one order better than before (I-1-1 in Table 2.1) thanks to the high spatial
resolution (0.4"). The nanoflare heating model predicts transient hot (>5 MK) and fast plasma flows(~100 km/s), that is attempted to be captured by the high-sensitivity spectroscopic observation (I-12 in Table 2.1). The high spatial resolution is crucial to clarify where and how nanoflare heating
occurs in association with the fine magnetic structures in the corona (I-1-3&4 in Table 2.1).
I-2: The wave heating hypothesis suggests that MHD waves, especially the Alfvén waves,propagate upwards along the magnetic field lines from the solar surface and dump their energy toheat the atmosphere. The spectroscopic capability of Solar-C_EUVST will allow precise
measurements of velocity fluctuations caused by transverse oscillation of magnetic fields (I-2-1 inTable 2.1). The seamless coverage of all temperature regimes is essential to catch the propagationof Alfvén waves from the chromosphere to the corona. The high spatial resolution is also critical toobserve fine magnetic structures interacting with waves and generating turbulent plasma motions,which is predicted by the theories of thermalization process (I-2-2 in Table 2.1). The plasma
diagnostics by Solar-C_EUVST will be complemented by high resolution magnetic fieldobservations from the ground-based telescopes, providing a complete picture of injection,
propagation, and dissipation of the waves (I-2-3 in Table 2.1). These observations provide
quantitative understanding of how much energy flux is carried by the Alfvén waves and where theenergy is dissipated.
I-3: Spicules are the fundamental constituents in the chromosphere and considered to be the key
player in channeling the chromosphere and corona by dynamic jets. Solar-C_EUVST will trace
thermal evolution of spicules from 2×104 K in the chromosphere to higher than 1 MK in the corona
by utilizing its seamless coverage over the whole temperature range as well as high spatial
resolution enough to discern each spicule. Thermodynamic properties obtained by the spectroscopic
observations enable to estimate energy and mass flux into the corona carried by spicules (I-3-1 inTable 2.1). Because spicules originate in the photosphere and lower chromosphere, their driving
mechanisms are investigated by joint observations with ground-based telescopes (I-3-2 in Table
2.1).
I-4: The solar wind is the plasma flowing along open field lines into the heliosphere. It is expected
to find clear evidence of acceleration by Alfvén waves as it flows away from the corona. The highsensitivity spectroscopic capability of Solar-C_EUVST aims to get profiles of velocity,temperature, and density along field lines from a faint source region near the surface to the upper
corona (I-4-2 in Table 2.1). The measurements yield the height dependence of the Alfvén wave
amplitude and its relationship with the plasma flow and temperature for different coronal structures
such as polar plumes and inter-plumes (I-4-2 in Table 2.1). It provides a clue to resolve thediscrepancy between theory and observation on how much the Alfvén wave dumps their energy todrive the flow.
II-1: Magnetic reconnection is one of the fundamental processes for converting magnetic energy
into the thermal and kinetic energy of the plasma. This process occurs much faster than is predictedby classical theory. Solar-C_EUVST will observe the dynamics of magnetic structures tounderstand the mechanisms that lead to fast magnetic reconnection in partially or fully ionizedplasmas. Solar-C_EUVST will test reconnection models (e.g., Sweet-Parker, Petschek, or plasmoidunstable type reconnection) by observing the velocity, temperature, density, ionization rate in andaround the shocks, and magnetic islands of the reconnection region (II-1-1 in Table 2.1). The highspatial and temporal resolution will allow us to study chromospheric evaporation produced by
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reconnection not only at coronal temperatures but also at low temperatures deep in the transition
region and chromosphere (II-1-2 in Table 2.1). This will determine how much energy is released
through magnetic reconnection. The observation of magnetic reconnection in the chromosphere by
Solar-C_EUVST will clarify the role of the fast reconnection process in weakly ionized plasma (II1-3 in Table 2.1).
II-2: Understanding of the accumulation and release of free magnetic energy is one of the last piecesof the "solar flare" puzzle. To identify the signatures of global energy buildup, Solar-C_EUVST
will perform long-term, wide-field monitoring of the AR corona and investigate the evolution of
bulk and turbulent (i.e. non-thermal) plasma velocities (II-2-1 in Table 2.1). By taking advantagesof its high spatial, temporal, and spectral resolution, Solar-C_EUVST will focus on the core of theAR and give us detailed information of the flow fields around the flare-triggering site (II-2-2 inTable 2.2). These observables are critically important not only for clarifying the physical conditions
and MHD instability modes of the flare eruptions but also for building the algorithms for flareprediction.
Some of the objectives above mentioned require measurements of magnetic fields, although nopolarimetric capability is included in the measurements of Solar-C_EUVST. Morphologicalpatterns of the intensity and velocity are only a tracer of the magnetic field structures. A mitigationis, however, to extensively coordinate its observations with ground-based observatories, especially
DKIST, which will provide measurements of photospheric and chromospheric magnetic fields withhigh spatial resolution, although the field of view is limited. These coordinated observations will beessential for evaluating the energy flux carried by waves from the chromosphere toward the corona.
3.3.

The compelling nature of the proposed mission concept and its relationship to past, current,
and future other investigations and missions.
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Before arguing the compelling nature of the proposed mission concept in section 3.3.2, current andfuture related missions are briefly described in section 3.3.1.
3.3.1 Past, current, and future
Japanese Solar Missions
Three solar physics missions have been launched by ISAS/JAXA in the past 40 years: Astro-A(Hinotori), Solar-A (Yohkoh) and Solar-B (Hinode). Yohkoh and Hinode have been International
Partnership Missions with the US (NASA) and the UK (PPARC, STFC), and Europe (ESA on
Hinode). All three missions have been highly successful and are providing a wealth of discoveries.
Based on these successes, we have established reliable and strong partnerships with the US andEurope. Solar-C_EUVST will build upon this rich tradition.
Launched in 1981, Hinotori provided the first hard X-ray images of solar flares. It clarified theimportance of hard X-ray imaging observations for understanding particle acceleration in the flares.
Yohkoh, launched in 1991, was devoted to observations of solar flares and the solar corona using a
combination of soft and hard X-ray imaging, and non-imaging spectroscopy. It is not an
exaggeration to say that the bulk of our knowledge of high-energy solar phenomena has come from
Yohkoh. For example, Yohkoh provided strong observational support for concluding that magnetic
reconnection plays a key role in the energy release process for flares.
Hinode, launched in 2006, observes photospheric vector magnetic fields, and their energetic anddynamical consequences in the corona. Hinode carries three main instrument packages: a Solar
10/100Optical Telescope (SOT), an X-Ray Telescope (XRT), and an Extreme-Ultraviolet Imaging
Spectrometer (EIS). The SOT has provided photospheric/chromospheric images and the spectropolarimetric data for measurements of the vector magnetic field in the photosphere, having 0.2"0.3" spatial resolution. The vector magnetograms produced by Hinode have the highest sensitivity
and spatial resolution ever obtained. The EIS has discovered extensive outflows at the edges ofactive regions that could be important contributors to the solar wind and is providing valuable
information on solar flares, coronal heating, and mass inputs from the chromosphere. The XRT has
provided new insights into hot plasma created by the dynamical events in solar flares and ARs.
Related Solar Missions Currently Running
SoHO : Solar and Heliospheric Observatory, launched in 1995, is an international project between
ESA and NASA to study the Sun from the interior to the outer corona, including the solar wind.
CDS (Coronal Diagnostic Spectrometer) and SUMER (Solar Ultraviolet Measurements of Emitted
Radiation) detected EUV and VUV emission lines in the chromosphere through the transition
region to the corona, providing diagnostic information on the solar atmosphere (flows, temperature,and density). Currently, only some of the instruments on SoHO are still providing data formonitoring the full Sun, and none of them has any overlap with the high-resolution instruments.
RHESSI : Reuven Ramaty High Energy Solar Spectroscopic Imager is a NASA SMEX mission
launched in 2002. RHESSI's primary mission is to explore the basic physics of particle accelerationand explosive energy release in solar flares. This is achieved through imaging spectroscopy in Xrays and gamma-rays to reveal the locations and spectra of the accelerated electrons and ions and of
the hottest flare plasma. Its limited spatial resolution does not allow it to address the issue ofmagnetic reconnection, which Solar-C_EUVST will challenge.
SDO : The Solar Dynamics Observatory, launched in 2010, carries three instruments. HMI(Helioseismic and Magnetic Imager) and AIA (Atmospheric Imaging Assembly) are instruments forfull-Sun observations. HMI is designed for measuring the magnetic field on the surface. AIA
images the chromosphere and the corona in multiple UV wavelengths to link changes in the outer
atmosphere to the surface. EVE (Extreme Ultraviolet Variability Experiment) measures the solar
EUV irradiance (i.e. no spatial resolution) with high spectral resolution, temporal cadence, andprecision. Although SDO does not provide any spatially resolved spectroscopic information, it will
supply context coronal images and surface magnetograms that have a high potential for synergy
with the high-resolution EUVST data.
IRIS : Interface Region Imaging Spectrograph is a NASA SMEX mission launched in 2013. IRIS
obtains UV spectra with high resolution in space (0.33-0.4") and time (4-8 s) focusing on thechromosphere and lower transition region. IRIS has very little coverage of lines formed in the hottransition region or corona, and thus cannot provide the observation that seamlessly cover the wholetemperature range of the outer solar atmosphere.
Heliospheric Missions in the Near Future
PSP : Parker Solar Probe (to be launched in summer 2018) is a mission to trace the flow of energy
through the inner heliosphere and to explore what accelerates the solar wind as well as energetic
particles by flying to within 9 solar radii from the solar surface (first close approach in 2024). PSP
will carry a suite of four instruments (three in-situ measuring instruments and one remote-sensing
instrument) designed to study magnetic fields, plasma and energetic particles. WISPR (Wide-field
Imager for Solar Probe) will image the solar wind structures as they pass the spacecraft. PSP willprobe the microphysics of the outer corona and solar wind in a unique manner, but will not observe
atmospheric conditions near the solar surface. Solar-C_EUVST is therefore a complement to PSP.
11/100SO : Solar Orbiter (to be launched in 2020) will provide the first-ever images of the Sun from an
out-of-ecliptic viewpoint (up to 25° of solar latitude during the nominal mission phase until 2027,
and up to 34° during the extended phase) by in-situ and remote-sensing instruments. Among the 6
remote-sensing instruments, SPICE (Spectral Imaging of the Coronal Environment) will obtain
UV/EUV spectra with 4” spatial resolution (equivalent to 1.3" at 1 AU at the closest approach of
0.3 AU) to detect plasma properties of the Sun’s on-disc corona. STIX (Spectrometer/Telescope for
Imaging X-rays) provides imaging spectroscopy of X-ray emission to measure timing, location, andspectra of accelerated electrons as well as high temperature plasmas associated with flares andmicroflares. Although SO does not measure imaging spectroscopic data with high spatial andtemporal resolution, it will give us observations from different view angles, which has a very highpotential for synergy by coordinating with Solar-C_EUVST (e.g., 3D tomography analysis).
Missions Under Study
MUSE : Multi-Slit Solar Explorer is a proposed NASA Small Explores (SMEX) mission currently
under phase A study for investigating the dynamics of the corona using novel spectral
imaging techniques. MUSE will obtain EUV spectra and images with the highest spatial (0.3") andtemporal (1-4 s) resolution in a large FOV, enhanced by multi-slit spectroscopy of a few selected
spectrum lines. Although MUSE has limited temperature coverage as well as diagnostic capability(e.g., temperature, density diagnostics), MUSE is a revolutionary type of coronal imager. It isreally T-07 in the NGSPM-SOT report and is very complementary to Solar-C_EUVST or T-09, notredundant to or directly competitive with it.
FOXSI : The Focusing Optics X-ray Solar Imager is a proposed NASA SMEX mission currently
under phase A study for investigating acceleration of energetic particles and the heating of the solar
corona to high temperatures. The FOXSI SMEX concept is based on the FOXSI sounding rocket
payload, which was flown in 2012 and 2014. A third launch is planned for 2018. FOXSI’s direct
imaging will provide a much higher dynamic range than previous imagers, such as RHESSI. FOXSI
would complement Solar-C_EUVST by observations of higher temperature plasma than isaccessible with EUV spectroscopy. FOXSI, however, would have significantly lower spatial
resolution (8") than Solar-C_EUVST and not provide information such as Doppler shifts andplasma densities.
3.3.2 The compelling nature of the proposed mission concept
The Solar-C_EUVST mission will generate spectroscopic observations with vastly improved spatial,
temporal, and temperature resolution. These improved capabilities will allow for two types ofobservations that have previously been impossible. First, the high spatial resolution of SolarC_EUVST will allow for very small scale structures to be observed over the full range oftemperatures in the solar atmosphere (see Figure 3.3). This will allow the plasma properties of thesestructures to be observed directly for the first time. Such spectroscopic observations will provide
diagnostics of Doppler shifts, non-thermal velocities, temperature, and density at a spatial resolution
and cadence currently only achievable by imaging instruments. This ability to observe 2D plasma
properties at high cadence for the complete atmosphere has the potential to be as transformationalas the first high-resolution imagers of the corona. Detailed comparisons of the performance ofSolar-C_EUVST and other spectroscopic missions will be presented in section 6.2.5.
1) Spatial resolution
Observations made with Hinode and other missions indicate that the spatial resolution required forobservations of elemental structures and processes is within our reach. This conclusion is based inpart on the Hinode/SOT observations that show that the typical chromospheric structures (such as
spicules, Pereira et al. 2012) are 300 km (0.4"). Consistent with this, studies of coronal rain,condensations of million-degree plasma that form in the corona when hot loops cool very suddenly,
12/100show widths of about 300 km (Antolin & Rouppe van der Voort 2012). Similarly, estimates of thefilling factor of EUV coronal loops derived by comparing the density and observed intensity
suggest loop widths of several hundred km (Hinode/EIS observation by Warren et al. 2008).
Finally, the modeling of evolving loops at 1 MK temperatures also suggests that they could becomposed of only a few strands, each several hundred km in width (Brooks et al. 2012). The SolarC_EUVST mission will attempt to resolve such elemental structures. It is important to note that
there is almost certainly a distribution of loop widths in the solar atmosphere, one that may extend
down to loops only a few tens of km wide. The accumulated evidence, however, suggests that there
is a strong peak in this distribution near several hundred km (e.g., Aschwanden & Peter 2017).
2) Temporal resolution
We need measurements of both thefast time-scale MHD perturbations ofstructures due to waves, as well as
measurements of the slowly varying
magnetic structures subject to quasistatic deformation. Such measurements
will allow us to determine therelationship between the field topology
and atmospheric heating, as well as theenergy flux carried by waves and their
dissipation. As another example of thenecessity of high time cadence, let us
consider observations of extremely
high temperature emission (>5 MK).
High temperature plasma is produced Figure 3.3: Characteristics of Solar-C_EVUST in terms of theimmediately after an impulsive energy temperature and spatial coverage, as compared with the currently
input (during a flare, CME, jet, or running missions.
nanoflare). Because the lifetime ofsuch high temperature emission can be very short (~ a few seconds), high temporal resolution
observations are critical. In order to achieve the required high temporal resolution, the effective areafor Solar-C_EUVST is ~ 10 times larger than previous instrumentations (see Figure 6.1).
3) Temperature coverage
Another unique approach of the Solar-C_EUVST mission is to perform photometric and Doppler
observations without temperature gaps from the chromosphere to the corona. Hinode is blind tomuch of the transition region and has no diagnostic capability in the chromosphere. It has become
clear from analyses of Hinode data that these observational gaps prevent complete understanding
not only of these layers, but also part of the corona and heliosphere, since the solar atmosphere is a
coupled system. Solar-C_EUVST will rectify this with an instrument that simultaneously observes
emission from the chromosphere, transition region, and corona at similar high spatial resolution (seeFigure 3.3). This is one of the most innovative and challenging properties of this instrument.
4.
4.1.

Launch constraints
Target launch date and rough schedule [18-1-1. 18-2-1]
J<=@;B H9DWRVaid

13/100Figure 4.1: Strawman schedule

Figure 4.1 is the overall strawman schedule
for the launch on February 2025 (JapaneseFiscal Year (JFY) 2024). This is the earliest
schedule. For maximizing the scientific returnby increasing the possibility of capturing
many large flares and by coordinating withobservations by SO (expected to be over 25°
helio-latitude around 2024) and the PSP(expected to have a close approach <10 solar
radius around 2025), the launch in JFY 2024
is preferable. JFY2026 is also acceptable,
based on the predicted number of flares(Figure 4.2). In the strawman schedule (Figure4.1), the time scale before System Definition
Review (SDR) is very tight for starting the
project from the beginning of 2020.
4.2.

Figure 4.2: Predicted number of large (>M-class) flares ineach year. Red columns give the number of large flares
actually observed by a similar spectrometer EIS. The solar
minimum is assumed in 2019, with shifting Hinode’s
statistics from 2006 to 2017.

Constraints from the launch vehicle (launch performance, envelope etc.) [13-2-1]
@; 9LF+67@;/j2Pg^fi]DC)An Epsilon vehicle after H3 synergy update (completed by 2021) is assumed. The launch
performance is the same as the performance of the 2nd flight (December 2016 for ERG), the detailsof which are described in Epsilon Rocket Users’ Manual (NC, March 2016).
Table 4.1: Accuracy (3σ) of the orbit installation. Orbit parameter
Configurations / Orbits
Perigee altitude [km]Apogee altitude [km]Inclination [deg]Without PBS
LEO±25 (250)±100 (500)±2.0 (31.0)
Elliptical orbit±25 (250)±2000 (30700)±2.0 (31.0)
With PBS
LEO±20 (500)±20 (500)±0.1 (30.5)
SSO±20 (500)±20 (500)±0.2 (97.4)

14/100For installation to a sun synchronous polar orbit (SSO), PBS (Post Boost Stage) is required toguarantee the accuracy of the installed orbit. The accuracy is listed in Table 4.1; for the SSO, +/-20
km in the altitude and +/- 0.2 degrees in inclination. The maximum weight for the SSO is given in
Figure 4.3(a). The Epsilon team (R. Yamashiro) suggests the case that the PBS has a propellent
mass of 160 kg; namely, the maximum weight is 590 kg for the height of 600-650 km and 560 kg
for 700 km. If more weight is needed for the payload, low earth orbit (LEO) shall be considered.
Figure 4.3(b) shows the envelope of the Epsilon rocket nose fairing, in which the proposed mission
should be accommodated.

Figure 4.3: (a) The payload weight for a sun synchronous polar orbit. (b) The nose fairing envelope.

5.
5.1.

Scientific investigations of the mission [3-2–1]
`YUbg: =M')h
Rationale for the investigations [3-2-2][1-3-1]
`YUbg: =M')hF

To achieve the science objectives of the mission which are shown in section 2 (Table 2.1), there areseveral possible instruments that can be considered. To discuss the rationale for the investigations,
we assume a slit-scanning type imaging spectrometer to clarify the points. Note that this type ofspectrometer will be chosen by a trade off study described in section 6.2. In this section, Table 2.1
is referred for the number given in science sub-objectives/observing tasks, such as I-1/I-1-1,
respectively.
I-1: Quantify the Contribution of Nanoflares to Coronal Heating (Table 5.1)
One of the main science objectives of Solar_C-EUVST will be to quantify the contribution of
nanoflares to coronal heating. Since the Skylab era (e.g., Rosner et al. 1978) it has been known thatthe coronal loops in the core of an AR have temperatures of 3-4 MK. These bright, hot AR loops
will be the primary target for this science objective. As is illustrated in Figure 5.1(a), these AR coreloops can be observed in emission lines from Fe XVIII. This figure, which shows an observation
with the AIA 94 Å channel, illustrates an imaging observation of the Fe XVIII 93.94 Å line. SolarC_EUVST will observe the longer wavelength Fe XVIII 974.86 Å line, which has the highest
photon flux per unit emission measure of any Fe XVIII transition. The object is to search for fainthigh temperature emission predicted by nanoflares in all the AR loops.

15/100Figure 5.1: (a) SDO/AIA routinely observes transient heating events in Fe XVIII. Solar-C_EUVST will observe
them with much higher spatial resolution and cadence and also provide plasma diagnostics such as temperature,density, Doppler shift, and non-thermal broadening, the quantities that should be compared with theories. (b) SolarC_EUVST will be able to observe the evaporative upflows associated with small-scale heating events predicted bynanoflare models. Here a simulated footpoint profile (blue) and loop-averaged profile (red) are shown. (c) Highspatial resolution images from Hi-C show braided loops in Fe XII. These loops are associated with the formation ofhigh temperature plasma. With its high spatial resolution Solar-C_EUVST will observe such events over a very wide
range of temperatures. (d) Hinode/SOT observations of the magnetic field at the footpoints of the Fe XVIII loops inpanel (a). Similar coordinated observations with high resolution ground-based observatories such as DKIST willallow us to reveal the photospheric changes that drive nanoflares to be identified. Schematic illustration of nanoflareon the top of figure is made by ISAS/JAXA.

The improved spatial resolution and sensitivity of Solar-C_EUVST will allow events with smaller
energies to be studied. The most compelling evidence for this comes from the Hi-C sounding rocket
experiment, 5 minutes observations of the footpoints of high temperature loops. Hi-C took very
high spatial resolution (0.3-0.4") images at a cadence of about 5.5 s. Testa et al. (2013) showed thatwith the increased spatial resolution and cadence of Hi-C, rapid variations in footpoint intensity
could be detected. The estimated energies are ~1024 erg for these events. Furthermore, these events
could not be detected at the lower spatial resolution and cadence of SDO/AIA.
The first observing task (I-1-1) will be to quantitatively measure the energy of heating events withenergies in the nanoflare range (~1024-1027 erg) envisioned by Parker (1972). These events can only
be detected at the higher spatial resolution (0.3—0.4") and cadence (5.5 s). Solar_C-EUVST will beable to follow nanoflare’s evolution in the core of the AR from high temperatures to lowtemperatures (wide temperature coverage) with high temporal resolution. The spectroscopic
observations will also provide electron densities, further improving the estimate of event energies.
16/100

These detected events will be used to form a statistical distribution of event energies, which will becompared against the energy requirements of the active corona (see Shimizu 1995 for the seminal
paper on this type of analysis).
The nanoflare concept predicts that heating events should be associated with chromosphericevaporation similar to that observed in large flares. As with large flares, the complicating factor isthe superposition of emission along the line of sight. As modeled by Patsourakos & Klimchuk(2006), spectral line profiles that have enhancements in the blue wing are expected from nanoflareheating, not completely shifted line profiles. This is illustrated in Figure 5.1(b), which shows thatthe extent of this enhancement is determined by the amount of energy released and the spatial scaleof the energy release. The second observing task (I-1-2) will be to observe these evaporative
upflows and compare the observations to model predictions. This will be achieved using sit-andstare observations of Fe XVIII at a cadence of 5 s or better.
The third observing task (I-1-3) will be to search for the field line braiding envisioned in the Parker
model. As is illustrated in Figure 5.1(c), in the Hi-C observations these twisting field lines wereassociated with the formation of high temperature plasma (Cirtain et al. 2013). Solar-C_EUVST
will observe this process over a much wider range of temperatures and also provide the full range ofplasma diagnostics that is needed for comparing with theory. This goal will be achieved using the
combination of slow (a few minutes) and full spatial resolution (0.4") scans and fast (~ a few 10 s)
middle resolution (0.8") scans in a wide selection of spectral lines over a small field of view.
The fourth observing task (I-1-4) requires the same observational parameters as (I-1-3). The
important requirement is coordinated observations of the photosphere (Figure 5.1(d)) and low
chromosphere by a ground-based telescope.
Table 5.1: Summary of requirements for science objective I-1
Item

Requirement
I-1-2
I-1-3
I-1-4
0.8-10(5-10 for nanofare signatures)

Rationale

I-1-1
Temperature
coverage (MK)Spatial resolution(arcsec)

0.4

0.8

Temporalresolution(s/step)

5
/0.4"

1.5
/0.8"

FOV
(arcsec2)

0.4
×
100

0.8
×
280

Velocity resolution(km/s)
Observation
duration & total
days

0.4
[high cadence:
0.8]
5/ 0.4"[high cadence:
1.5/0.8"]

20×100[high cadence:
20×100]Vd ~ 2
Vnt ~ 5
~1 hour
15 days each

To observe nanoflare heating-cooling cycles. Nanoflaresignatures are detected with high temperature lines.
0.4" for resolving nanoflare events and subarcsec braiding
structures and 0.8" for resolving fine structures at footpointsof coronal loops.
I-1-1: Similar cadence as Hi-C observation (5.5 s).
I-1-2, I-1-3: The exposure time is shorter than the spatial
scale (0.8" ~ 600 km) / sound speed 300 km/s, ~2 s.
I-1-1: The slit on the core of an active region.
I-1-2: Capture both of the footpoints of AR loops.
I-1-3: Capture the entire or a part of coronal loops.
Enough to resolve plasma flows (Doppler and turbulent
components) caused by tiny energy releases.
Longer than the evaporation-condensation cycle.
Repeated measurements for statistical studies

For precise co-alignment with ground-based data.
Simultaneous slit jaw imaging to know where the slit islocated on solar images.
1": 730 km, Vd: Doppler velocity, Vnt: nonthermal velocity (e.g., turbulence)
Co-alignment

n/a


√

17/100I-2. Quantify the Contribution of Wave Dissipation to Coronal Heating (Table 5.2)
The magnetized solar atmosphere is filled with a rich assortment of waves that are an importantsource of energy flux. The exact amount of energy flux they transport and the fraction of the energy
dissipated into various structures still need to be determined. In particular, Alfvén waves have long
been proposed as a candidate mechanism to transport magneto-convective energy from theconvection zone into the solar atmosphere where they can heat plasma to coronal temperatures andaccelerate solar wind through dissipation of their non-thermal energy (e.g. Suzuki et al. 2005, vanBallegooijen et al. 2011). Because Alfvén waves are subject to wave reflection (van Ballegooijen etal. 2011), as the Alfvén speed increases from about 10 km/s in the chromosphere to more than 1000
km/s in the corona, especially in the transition region, only a fraction of the wave energy can betransmitted into the corona. It is thus critically important to clarify how the velocity fluctuation
propagates from the chromosphere to the corona through the transition region. This is the firstobserving task (I-2-1) of the science objective that is attempted by simultaneous observations ofemission lines from 0.02 to 3 MK in the atmospheres. If Alfvén waves carry a significant amount of
the energy transmitted in the corona, 105 erg/cm2/s, a velocity amplitude of Alfvénic waves is
expected to be of the order of 10 km/s in the corona. Recent imaging observations have detectedsignatures of transverse oscillations in chromospheric structures (Figure 5.2(a)), suggesting thepotential of waves in powering the solar atmosphere (De Pontieu et al. 2007a; Okamoto et al. 2007;
Jess et al. 2009; Okamoto & De Pontieu 2011). Alfvénic waves have been observed in the corona
by a coronagraph observation (e.g. Tomczyk et al. 2007), but the velocity amplitude is smaller than
1 km/s, giving an estimate of flux which is much smaller than is required for coronal heating. Thismay be largely underestimated due to the integration of signals over unresolved features in thecoronagraph observation. Thus, spatially resolving coronal structures with Solar-C_EUVST isessential to accurately determine the velocity amplitude of waves and to evaluate the magnitude of
their energy flux.

Figure 5.2: (a) Observation of a prominence thread structure indicating lateral and line-of-sight oscillation (Okamoto
et al. 2015), (b) synthesized line intensities from the chromosphere (104 K) and transition region (105 K) indicating
appearance in the hotter line as a result of thermalization of the Alfvén wave (Antolin et al. 2015), and (c)observation of velocity and magnetic fluctuation to explore a source of the Alfvén wave (Kanoh et al. 2016).

The theory for the wave that Alfvénic waves interact with density structures in the transition region
and corona, which causes nonlinear evolution of the velocity amplitude causing mode conversion
18/100and turbulence, leading to dissipation and heating (e.g. Okamoto et al. 2015, Antolin et al. 2015,
Matsumoto 2016) (Figure 5.2(b)). In the second task (I-2-2), Solar-C_EUVST will measure
spectroscopically how amplitudes and phases of velocity fluctuation vary and are correlated amongemission lines from the chromosphere to the corona. The observation is done along withobservations of the density structure with the sub‐arcsec resolution, needed to characterize thetransverse variation of the atmosphere structure. Solar-C_EUVST should scan a footpoint region ofa magnetic field line to achieve a temporal cadence shorter than 1 min with the FOV narrower than
10" to get high spatial and velocity resolution to yield key information on the wave propagation anddissipation processes. Another observing mode is to map a 200 Mm wide region in a time shorter
than the Alfvén speed (~1000 km/s in corona), i.e., less than 200 s in order to follow the wave as it
moves through the corona and see its influence on the upper corona. This will be achieved byrepeated raster scannings with a wide slit (1.6"). Because Solar-C_EUVST is sensitive only toatmospheric layers hotter than 0.02 MK, sources of the Alfvén waves in the photosphere andchromosphere are studied by coordination with ground-based telescopes, such as DKIST.
Table 5.2: Summary of requirements for science objective I-2
Item
I-2-1

Requirement
I-2-2
I-1-3
0.02-3

Rationale

Temperature
coverage (MK)Spatial resolution(arcsec)

0.4[high cadence: 1.6]

Temporal resolution(s/step)

2/0.4"[high cadence: 0.5/1.6"]

FOV(arcsec2)

10×100[high cadence: 280×280]

It is possible to cover the wavelength of Alfvén waves with aperiod shorter than 200 sec. 1000 km/s × 200 s ~ 200 Mm,
corresponding to 280".

Velocity resolution(km/s)
Observation duration& total days

Vd ~ 2For measuring transverse oscillations 10 km/s responsible forenergy transfer.
Cover more than 10 wave periods where the typical period is 3
to 5 minutes.
Repeated measurements for statistical studies
For precise co-alignment with ground-based data. Simultaneous
slit jaw imaging to know where the slit is located on solar
images.

Co-alignment~1 hour
15 days each
n/a


√

To observe velocity fluctuation which propagate from thechromosphere to the corona through the transition region.
For resolving thread-like structures through the chromosphere to
the corona, 0.4" is required. For achieving high S/N and hightemporal resolution, 1.6" is acceptable.
We make use of a fast scan to trace the propagation of Alfvén
waves whose speed is 1000 km/s in the corona.

1": 730 km, Vd: Doppler velocity

I-3. Understand the formation mechanism of spicules and quantify their contribution to coronalheating (Table 5.3)
Observations of spectral lines from the chromosphere, such as Ca II, have clearly revealed that thechromosphere is dominated by a multitude of thin (300 km, i.e., 0.4" width, see Figure 5.3),
dynamical, jet-like extrusions called spicules. They have a typical upflow speed of 50 to 200 km/s
and a lifetime of 20 s to 300 s (e.g. Pereira et al. 2012). The origin of spicules is still not understoodand the contribution of spicules to the million degree corona has not been determined. Spicules havebeen reported to produce broadening and asymmetries in coronal spectral line profiles observed byHinode/EIS, whose spatial and temporal resolution is 2" and 300 s, respectively (De Pontieu et al.
2009). Such observations have led to the hypothesis that spicules are an important contributor ofmass and energy to the corona, although the sizes and temporal evolution of spicules are well below
the resolution of current UV and EUV spectrometers such as Hinode/EIS. Spatial and temporal
resolution in the spectroscopic measurements is crucial to understand whether spicules provide asignificant contribution to heating at transition region or coronal temperatures. IRIS achieved
sufficient spatial resolution to resolve spicules (Pereira et al. 2014), but it provides temperature
19/100coverage only in the chromosphere and the low transition region. Although radiative MHD
simulations have some success in reproducing spicules, their role in coronal heating (Figure 5.3) isstill not clear. Solar-C_EUVST will allow us, for the first time, to trace individual spicules
seamlessly as they rise into the corona by monitoring the evolution of the line profiles formed atincreasingly higher temperatures. In the first task (I-3-1), repetitive scans for a 10" wide area areneeded with high spatial resolution (0.4") as well as high cadence (180 s < life time of spicules).
Although Solar-C_EUVST cannot observe the photosphere and low chromosphere, the second task
(I-3-2) will be achieved through coordination with ground-based telescopes such as DKIST.

1. Figure 5.3: Left) Sample spicule Ca II H image observed by Hinode/SOT (Pereira et al. 2012). The typical spicule
width can be estimated to ~0.4". Right) Spicule-like features in a radiative MHD simulation (Martínez-Sykora et al.
2017). We can clearly see that the cold spicules penetrate into the hot corona.

Table 5.3: Summary of requirements for science objective I-3
Item
Temperature
coverage (MK)Spatial resolution(arcsec)
Temporal resolution(s/step)
FOV(arcsec2)
Velocity resolution(km/s)
Observation duration& total days
Co-alignment

Requirement
I-3-1
I-3-2
0.02-3For observing the chromosphere, transition region, and corona simultaneously.

0.4

For resolving sub-arcsec spicule structures in the chromosphere and theircounterparts in transition region and corona.

2/0.4"For scanning a 10" wide area within 1 min. Shorter than the life time of
spicules (20~300 s).

10×100The entire structure of spicules and chromospheric jets should be observed.

Vd ~ 2

Enough to resolve upward speed of spicules.

~1 hour
15 days eachn/a√Rationale

Longer than the typical transient phenomena in chromosphere (e.g., jets).
Repeated measurements for statistical studies
For precise co-alignment with ground-based data. Simultaneous slit jaw
imaging to know where the slit is located on solar images.

1": 730 km, Vd: Doppler velocity

I-4. Understand the Source Regions and the Acceleration Mechanism of the Solar Wind (Table
5.4)
The source regions of the fast and slow solar wind are still under debate. Regarding the fast solar
wind sources, polar plumes (structures are about 10" wide at the base of the corona that have been
tracked while they expand super-radially up to about 10 solar radii and beyond, and their lifetimes
are about 10 hour, Wilhelm et al. 2011) and inter-plume regions, which extend from coronal holes
into the high corona and probably link to the solar wind, have been proposed as sources. Since
smaller outflow speeds are detected at the base of polar plumes (0.7~0.9 MK), inter-plumes (1~1.3
MK) have been considered to be the fast wind source. However, contradictory results have beenobtained from off-limb observations, with inter-plume regions showing larger outflow velocities insome studies (e.g., Teriaca et al. 2003) and plumes themselves exhibiting higher outflow velocities
20/100in others (Gabriel et al. 2003). For the slow solar wind, Hinode/XRT and Hinode/EIS have revealed
the presence of persistent flows (V ~ a few 10 km/s, T ~ a few MK) at the edges of ARs (e.g.,
Sakao et al. 2007) (size of AR ~ 280×280 arcsec2), which show the same compositional signatureas the slow wind observed at Earth (Brooks & Warren 2011). On the other hand, streamers in thehigher corona have also been proposed as one of the slow solar wind sources as a result of far offlimb observations (e.g., Raymond et al. 1997). So far, compelling observational evidence thatdefinitively identifies the source regions of the solar wind is still controversial, but it could be that a
combination of lower atmospheric sources (such as AR outflows) and higher altitude sources (suchas streamers) together provide the total mass flux of the slow wind (Brooks et al. 2015, Figure5.4a).

Figure 5.4: (a) Left) Plasma properties in solar wind source regions and their relationship with magnetic fields(Brooks et al. 2015). Left) full-Sun AIA 193Å composite intensity image (blue) with solar wind source regions (redand green), which is identified by abundance. Center) Overlay of magnetic field lines from the potential field
extrapolation calculation on the Fe XIII 202.044  Å Doppler velocity map from Hinode/EIS. Right) Percentage of
total mass flux as a function of magnetic field line-starting latitude (red). The vertical blue lines on the left and rightside mark the latitude of 11° and 40°. (b) Signatures of Alfvén waves in plumes and inter-plumes. Left) Plume andinter-plume observed by multiple telescopes (Teriaca et al. 2003). Center) Slit positions of Hinode/EIS for theobservation of the right panel overlaid on the SoHO/EIT 195 Å image. Right) off-limb observation of the nonthermal velocity (Vnt) from the strongest observed lines (symbols). The dashed line gives the predicted trend,
proportional to the electron density n as n-1/4, for undamped waves. The observed behavior of Vnt may be considered
as a damping signature of Alfvén waves (Hahn and Savin 2013).

Therefore, the first observing task (I-4-1) for the solar wind study is to diagnose the plasma
properties in these regions and clarify the source regions, their contributions, and the relationships
with magnetic structures. One difficulty in analyzing the source regions of the solar wind is that alonger temporal cadence (~hours) is needed to observe faint structures (such as coronal holes andAR boundaries). The throughput improvement of at least one order of magnitude by SolarC_EUVST will allow us to approach the timescales of their evolution (a few to several minutes).
Also, rolling the spacecraft by 90° will greatly reduce the scanning time for the equatorial limb
observations. Regarding the relationship between the solar wind source region and magnetic field
21/100(I-4-1), interchange reconnection between closed and open flux is widely thought to be the driverof the AR outflows. However, Hinode/EIS observations lack the spatial resolution to observe
small-scale changes low down in the atmosphere at the outflow/AR boundary, and they also lackthe temporal resolution to observe large-scale morphological changes that could result fromreconnection at higher altitudes. The improvement in spatial and temporal resolution of SolarC_EUVST will play a critical role in establishing which physical mechanisms are responsible forthe formation and acceleration of the slow and fast wind.
To understand the acceleration mechanisms of both fast and slow solar wind, the wave dissipation
rate in the corona is an important piece of information (e.g., Cranmer 2002). Therefore, the second
observing task for the solar wind (I-4-2) is detecting signatures of coronal Alfvén waves in plumeand inter-plume regions and measuring their energy fluxes with height. Hahn and Savin (2013)
discussed a measurement of the energy carried and dissipated by Alfvén waves in a polar coronalhole using line width measurements from Hinode/EIS (Figure 5.4(b), typically Vnt ~ 40 km/s,
which represents the Alfvén wave fluctuation amplitude). They found signatures of wave
dissipation from 1.15 to 1.5 solar radii, which is consistent with previous observational estimatesfrom SoHO/SUMER and SoHO/CDS (e.g., Banerjee et al. 1998). This dissipated energy flux
possibly accelerates and heats the solar wind. Dolla & Solomon (2008) have reported similar trends
in line widths observed off-limb, but they interpreted them in a different way. According to theiranalysis, instrumental stray light is significant for off-limb observations, and the effect needs to betaken into account when considering different heating mechanisms. To clarify the acceleration
mechanisms for the solar wind, we not only need higher sensitivity spectroscopic observations to
overcome the low sensitivity in faint and off-limb regions, but we also need to minimize the impact
of stray light by using highly polished optical surfaces (minimal micro-roughness) and by avoiding
entrance filters (that produce diffracted radiation). The improvement in spectral resolution (ex. Vnt~ 4 km/s) and reducing the effect of stray light of Solar-C_EUVST will allow us to investigate
more accurate energy flux for accelerating the solar wind.
Table 5.4: Summary of requirements for science objective I-4
Item
Temperature
coverage (MK)Requirement
I-4-1
I-4-2

Rationale

Plumes (0.7~0.9 MK) and inter-plumes (1~1.3 MK) in coronalholes and the edges of ARs (~3 MK). Chemical abundance
diagnostic capability.
Spatial resolution
For distinguishing fine structures in solar wind source regions,(arcsec)
0.8 arecsec
e.g. polar plumes and inter-plumes in polar coronal holes (width~ 10 arcsec).
I-4-1: About or less than 1 hour scanning large FOV (280"Temporal resolution
×280") is required, which is typical timescale of AR evolution.
(s/step)
1.5 s/ 0.8"
~5 min / 0.8" I-4-2: Shorter than the typical transient phenomena in polar
coronal hole (e.g., jets).
FOV
Off limb up I-4-1: The entire AR should be covered.
(arcsec2)
to 1.4 solar
280×280
I-4-2: Higher than 1.2 solar radii is required to quantify theradiiwave dumping with height.
Velocity resolution
Vd ~2 km/s
To observe a few 10 km/s velocity signal. To observe height(km/s)
Vnt ~ 4 km/s
dependence for the non-thermal velocity (typically ~40 km/s).
Observation duration~1 hourLonger than the typical transient phenomena, such as polar& total days
15 days eachcoronal hole jets.
Repeated measurements for statistical studies
For co-alignment precisely with ground-based data.
Co-alignment√n/aSimultaneous slit jaw imaging to know where the slit is locatedon solar images.
1": 730 km, Vd: Doppler velocity, Vnt: nonthermal velocity (e.g., turbulence)
0.7~3 MK

22/100II-1: Understanding of the Processes in Fast Magnetic Reconnection (Table 5.5)
The magnetic reconnection rate, how fast magnetic flux is swept through the reconnection region,
predicted by standard MHD theory is much too slow to explain the rapid evolution of events
observed in the solar corona. Although several theoretical models have been proposed to solve thisproblem (e.g., Petschek reconnection model Figure 5.5C, Plasmoid-unstable reconnection model
Figure 5.5D), the physics of fast reconnection is still not clear. The solar atmosphere is one of thebest space laboratories to answer this question, and observing the reconnection region (Figure 3.2A)
is crucial for tackling the problem (observing task: II-1-1). Figure 5.5(A&B) show imaging
observations of two flares by SDO/AIA. These images show a reconnecting current sheet structure
without islands (Warren et al. 2018) and with islands (Takasao et al. 2012). However, thereconnecting current sheet cannot be seen in most flares (e.g., Figure 3.2C). Inflows (~ a few 10km/s, Hara et al. 2011), shocks, and high-velocity jets that are expected inside the reconnection
region are also not observable in most cases. The reconnection region is generally very dark while
the post-flare loops are very bright. Therefore, we need high-throughput and low scatter-lightobservation to diagnose the reconnection region. The typical thickness of the current sheet andislands is a few arcsec, while the length is ~20". These islands and current sheets are widely
observed from 171 (0.8 MK) to 195 Å (~15 MK). Thus, wide temperature coverage (widewavelength coverage) is also important for understanding the dynamics of reconnecting current
sheets and islands. Recently several papers claim that the reconnection region may not be inionization equilibrium (e.g., Imada et al. 2011). To diagnose the non-equilibrium ionization plasma,
we need a series of spectral lines from several different atomic species, for example Ca and Fe,
which have different ionization timescales. Therefore, wide wavelength coverage is also very
important. Typical speeds of island motions are a few hundred km/s, requiring a cadence of a few
tens of seconds (15 Mm / 1000 km/s ~ 15 sec) to diagnose the reconnection region. Also, rolling thespacecraft by 90° will greatly reduce the scanning time for the limb flare region.

Figure 5.5: SDO/AIA observations showing a reconnecting current sheet structure (A) without islands (Warren et al.
2018) and (B) with islands (Takasao et al. 2012). Numerical simulations of magnetic reconnection (C) Petschek
reconnection (Yokoyama & Shibata 1997), (D) plasmoid-unstable reconnection (Shibayama et al. 2015).

23/100During a solar flare, thermal conduction and high-energy electrons progress down along themagnetic fields of loops and collide with the chromosphere, producing evaporation of hot plasma
into the corona (e.g., Fisher et al. 1985). The chromospheric evaporation generally includes
information concerning the energy conversion rate during magnetic reconnection (observing task:
II-1-2). To understand the chromospheric evaporation, observations over a wide temperature range
with high spatial and temporal resolution are needed since the kernels of chromospheric evaporation
are small (~1") and evolve on timescales of sub-seconds (105 km / 105 km/s = 1 sec; non-thermal
electron precipitation).
One of the most important Hinode findings is that magnetic reconnection events frequently occur
even in the solar chromosphere (Shibata et al. 2007). Then, the key question is whether themagnetic reconnection in the chromosphere is really fast, independent of resistivity, or not. Ourcurrent guess, which is based on the Hinode observations, is “yes”. Hinode Ca II H movies,
however, do not allow us to evaluate the reconnection rate (observing task: II-1-3). The timecadence of the Ca II H observations is fast enough, but quantitative measurements of e.g. flows,density, and temperature are needed to determine the reconnection rate. The typical parameters of
the reconnection in the chromosphere/transition region are as follows: current sheet thickness ~1",
length ~10", Alfvén velocity 100 km/s, timescale 10~100s (10" / 100 km/s), temperature 104~106 K.
Table 5.5: Summary of requirements for science objective II-1
Item
II-1-1

Temperature
coverage (MK)0.8-30Requirement
II-1-2

0.02-15

Spatial resolution(arcsec)

Temporalresolution(s/step)

FOV(arcsec2)

Rationale
II-1-30.02-30.8

0.5/0.8"

20×2800.8×280

1.5/0.8"

20×280II-1-1: To observe the magnetic reconnection region wherecool plasma (~1 MK) is heated up to 15 MK in the corona.
Non-equilibrium ionization diagnostic capability is also
important.
II-1-2: To observe chromospheric evaporation plasma withtemperatures from 0.02 (~chromosphere temperature) to 15
MK (flare temperature).
II-1-3: To observe the magnetic reconnection region wherecool plasma (~0.02 MK) is heated up to 3 MK in thechromosphere.
II-1-1: For resolving magnetic islands and the reconnecting
current sheet during the flare (~1").
II-1-2: For resolving kernels of chromospheric evaporation(~1").
II-1-3: For resolving reconnecting current sheet in thechromosphere (~1").
II-1-1:The scanning speed should be faster than the Alfvén
speed 1000 km/s in corona, i.e. 1 sec /step in 1" step.
II-1-2: Evolution timescale for the kernels of chromosphericevaporation (0.5 s).
II-1-3: The scanning speed should be faster than the Alfvén
speed 100 km/s in chromosphere, i.e. 7.5 s / 1" step.
II-1-1: The entire reconnection region in the corona should beobserved.
II-1-2: To observe flare ribbons evolution.
II-1-3: The entire reconnection region in the chromosphereshould be observed.
A few km/s to observe the reconnection related flows, such asreconnection inflow ~20 km/s.

Velocity
resolution
Vd ~2 km/s
(km/s)
Observationa few days
~1 hour A few days to catch at least one flare.
duration & total
2 months each around
15 days
To observe >30 events M class flare in total (Figure 4.2).
days
solar maximum
Repeated measurements for statistical studies.
1": 730 km, Vd: Doppler velocity, Vnt: nonthermal velocity (e.g., turbulence)

24/100II-2. Identify the Signatures of Global Energy Buildup and the Local Triggering of the Flare andEruption (Table 5.6)
Flares and CMEs, especially stronger ones, are known to occur in ARs. In the framework of thestandard flare model (Figure 3.2, e.g. Priest & Forbes 2002; Shibata & Magara 2011), a magnetic
flux rope is created through the large-scale evolution of an AR, and free magnetic energy (nonpotentiality) is accumulated in the atmosphere. The entire system becomes unstable and eventually
the flux rope erupts, perhaps triggered by small-scale processes such as flux emergence andmagnetic reconnection with the large-scale fields. Therefore, in order to understand thedestabilization and eruption, we need to reveal (1) the long-term evolution of entire ARs, as well as(2) the small-scale dynamics of the flare-triggering processes.
Long-term monitoring with a large enough field of view is essential for examining the signature ofenergy buildup in the corona from the early stage of AR evolution. Using Hinode/EIS, Harra et al.
(2009) detected an increase of 40-50 km/s in the non-thermal line widths of the Fe XII (195 Å) line
a few days before the flare onset. Imada et al. (2014) also measured an upflow of 10-30 km/s one
day before the flare and a loop expansion a few hours before (see Figure 5.6), both in the Fe XII(195 Å) line. However, the physical origin of these pre-flare phenomena and their relationship withthe global energy buildup remain unclear. To this end, Solar-C_EUVST will conduct monitoring of
entire ARs (field of view of e.g. 280"×280") for their full disk passage (up to 2 weeks) and, throughthe investigation of multiple spectral lines and comparison with photospheric evolution (observing
task: II-2-1), it will reveal how and where these upflows originate. To resolve the Doppler flows in
the lower atmosphere (e.g. chromosphere), where the sound speed is a few 10 km/s, a velocity
resolution of ~2 km/s is needed.
Observations have revealed that flares are often associated with small-scale, newly-emerging
magnetic flux and its trigger interaction with pre-existing fields (e.g. Feynman & Martin 1995;
Wang & Sheeley 1999). Photospheric observations (e.g. Toriumi et al. 2013; Bamba et al. 2013)
show that the triggering systems have spatial scales of 5"-20" (Bamba & Kusano 2018), and often
exhibit pre-flare brightenings in the chromosphere that probably indicate local reconnection withthe pre-existing fields. Since the dynamics of the trigger fields is determined by the Alfvén speed (afew 10 km/s), the evolution time scales should be less than 10 minutes. To quantify the trigger
dynamics, Solar-C_EUVST will perform repeated scanning around the polarity inversion lines withmedium spatial and temporal resolution (~1" and ~1 minute, respectively), simultaneously from thechromosphere to the upper atmosphere, and measure the flow fields (observing task:II-2-2). We will
identify the MHD instability modes by directly comparing the observational results with magnetic
measurements in the photosphere and numerical simulations. Understanding the flare trigger iscrucial for realizing flare predictions.

Figure 5.6: (Left) Preflare AR upflow observed by Hinode/EIS (Imada et al. 2014). Blue and red in the Dopplergram
indicate upflow and downflow, respectively. (Right) Local flare-triggering fields with small-scale reconnection(arrows) obtained by Hinode/SOT (Toriumi et al. 2013).

25/100Table 5.6: Summary of requirements for science objective II-2
Item

Requirement
II-2-1II-2-2

Rationale

II-2-1: A broad temperature range should be covered to monitor
pre-flare phenomena.
II-2-2: Lower temperature ranges should be covered to reveal
the dynamics of the lower atmosphere.
II-2-1: Resolution should be better than or similar to AIA
Spatial resolution
1.6
0.8(~1.6") is needed.
(arcsec)
II-2-2: High resolution observations should be conducted toresolve the flare-triggering fields (~1").
II-2-1: About or less than 1 hour scanning large FOV(280"x280") is required, which is the typical timescale of active
region evolution. Constant, continuous monitoring is crucial for
Temporal resolution
0.5/1.6"
1.5/0.8"
this observation.
(s/step)
II-2-2: Less than 10 minutes for the monitoring of a small FOV(20"×280") to track the evolution of the flare triggering region,and preferably about or less than 1 minute during the flaring
phase.
FOV
II-2-1: The entire AR should be covered.
(arcsec2)
280×280
20×280II-2-2: To observe flare triggering region.
II-2-1: For the coronal measurements, a few km/s for the
Doppler velocity and ~10 km/s for the turbulence velocity are
Velocity resolution
required to resolve, for example, the pre-flare dynamics andVd ~ 2
Vd ~ 2(km/s)
non-thermal broadening.
Vnt ~ 5II-2-2: For the chromospheric measurements, ~2 km/s for the
Doppler velocity is desirable so that the Alfvén and acoustic
velocities (a few 10 km/s) are resolved.
Observation durationa few days
A few days are requested to catch at least one flare event (up to& total days
1 month each2 weeks for the whole disk passage).
To observe >15 events M class flare in total (Figure 4.2).
For precise co-alignment with ground-based data. SimultaneousCo-alignment
n/aslit jaw imaging to know where the slit is located on solar√
images.
1": 730 km, Vd: Doppler velocity, Vnt: nonthermal velocity (e.g., turbulence)
Temperature
coverage (MK)0.02-100.02-1

26/1006.

Instrumentation of the mission [5-2–1]<;:=>0-8

6.1.

One sentence description of the technology to realize the investigations [5-2-1]
-8 %??329 -8 'A 0 #-8,1(&
The proposed mission carries a UV imaging spectroscopic telescope that has broad passbands in 170– 1280 Å with high spatial resolution of 0.4″, high throughput that achieve a high temporal resolutionof 0.5 s, and high spectral resolution to resolve velocity fluctuations of 2 km/s.
~

~

6.2.

~

Comparisons of the selected technology with other technologies [5-2-2]
-8 %??329 -8? 4.*/@5 @+76$
15!

In this subsection, first we compare possible technologies for the coronal observations from theviewpoint of observation parameters, and determine the best technology to achieve our scienceobjectives. Second, we examine wavelength bands and spectral lines for the selected technology.
Third we examine the radiometric performance of the chosen wavelength bands to confirm thecapability of plasma diagnostics for our science objectives. Fourth, we confirm the photometricaccuracy for the selected technology whether it satisfies our scientific requirements. After that wecompare the other missions that have similar technologies to that of the proposed mission.
6.2.1 Comparison of technologies for coronal observations
Table 6.1: Comparison of possible technologies. Yellow filled cells show the parameters that satisfy our scientificrequirements.
UV imaging UV imaging
SXR imaging SXR imaging SXRVisiblespectroscopy
spectroscopy spectroscopy imaging
spectroscopy(selected(micro(CMOS)
withtechnology)
calorimeter)
coronagraphTemporal
0.1-1
~1
<1
<10
<10
0.25
resolution(exposure) [s]Spatial
0.4
1.6 (AIA)
~1
1
1
~10
0.2 (Hi-C)
resolution [″]Spectral
~5000
N/A
~1000
5
N/A
~10000resolution[l/dl]
Temperature
20k – 15M
6k – 15M
6M - 50M> 5M
1.3M – 32M 10k, 1.5-2.5M
coverage [K]Temperature
0.05-0.1
0.2-0.8
0.1
0.1
0.2
N/A
resolution[d(logT[K])]full disk (AIA) ~1000<500
Full disk>2.8R_sun,
Field of view [″] 300
400 (Hi-C)
only off-limb
References:
UV imaging: SDO/AIA (Lemen et al. 2012), Hi-C (Cirtain et al. 2013)
SXR imaging spectroscopy (microcalorimeter): Prototype TES X-ray microcalorimeter array (Christe et al. 2016)
SXR imaging spectroscopy (photon counting with CMOS): PhoENiX (Sakao et al. 2014)
SXR imaging: Hinode/XRT (Golub et al. 2007)
Visible spectroscopy with coronagraph: NSO/COMP (Tomczyk et al. 2008)

27/100

The proposed technology should be the best method to obtain the data described in Tables 5.1 – 5.5.
Table 6.1 compares typical observation parameters obtained by different technologies for coronalobservations. All technologies in the table satisfy our scientific requirement of the temporal
resolution, and especially UV and soft x-ray (SXR) technologies satisfy the requirement of the fieldof view. UV technologies have a better spatial resolution, and meet the science objectives to observea wide temperature range from the upper chromosphere to the corona seamlessly and simultaneously.
UV spectroscopic and visible coronagraphic technologies have a capability to resolve the required
velocity structures. Therefore, we conclude that UV spectroscopic technologies are the best for thescience objectives described in the previous sections.
6.2.2 Temperature coverage of UV lines and spectral line selection
The observations must cover all relevant temperatures and structures to reveal the energy flows andenergy release signatures throughout the atmosphere. Tracking structures and flows through different
temperature regimes at 0.4″ resolution requires strict co-alignment of all spectral bands. Thus, aunique optical path to the grating assembly is a fundamental requirement and design driver. It isequally important to explore solar plasma at high temperatures above 5~MK. If the corona is heated
impulsively by numerous small-scale impulsive events, they are predicted to produce hot plasma
above 5~MK. Spectral lines from ions such as Fe~XVIII (7.1~MK) or ions formed at higher
temperatures can constrain the nanoflare heating scenario. Highly ionized iron ions, i.e., Fe~XVIII
through Fe~XXIV, also allow us to explore the temperature evolution for out of ionization equilibrium
plasma and therefore give new insights into the reconnection region.
UV chromospheric lines of the hydrogen Lyman series (20~kK) are located between 925 – 1085~Å.
Strong chromospheric lines such as N~I and the strongest hydrogen line, Ly-alpha 1216~Å, aredistributed between 1115 – 1275~Å. To observe continuously through the upper chromosphere to thelower transition region, O~II – O~IV (0.05 – 0.16~MK) lines located between 690 – 850~Å are crucial.
Lines formed between the upper transition region and the lower corona such as O V (0.25~MK) and
Ne~V – Ne~VII (0.32 – 0.63~MK) are located between 463 – 637~Å. Intense lines formed in the lower
corona are located between 170 – 215 Å, which are confirmed by the Hinode/EIS observations. Hightemperature AR lines such as Fe~XVIII 975 Å (7.1 MK) were a missing window of the Hinode/EIS
observations. Such observations can beachieved at the same time by observing
the passband 925 – 1085~Å. Hightemperature coronal emission from
Fe~XX – Fe~XXII is also located between
690 – 850~Å. Therefore, the required
temperature coverage resulted in the
careful selection of the six spectral
passbands, which consist of one shortwavelength (SW) band and five longwavelength (LW) bands as shown in
Figure 6.1. Note that two of the LW bands Figure 6.1: Selected spectral passbands and temperature
coverages. Strong lines are marked as diamond symbols.
are taken by second order diffraction.
The key requirement for the instrument is to make measurements with spectral lines formed over theentire solar atmosphere, i.e., at chromospheric, transition region, corona and flare temperatures withhigh spatial and temporal resolutions. This requires bright lines at all temperatures and has led to theselection of the six spectral bands. The brightest spectral lines in the passband are marked in Figure
6.1 and listed in Table 6.2. Note that a vastly larger number of fainter lines observed by EIS andSUMER are not marked but locate in the filled regions.
28/100Table 6.2: List of the brightest lines to be observed with Solar-C_EUVST. Green, blue, pink, yellow, andlight blue show lines formed in the chromosphere, upper transition region, lower transition region, corona,and flaring region.
Radiance [erg/s×cm2×sr]
ChannelIon
Log T [K]Flare
Wavelength [Å](order)QS
AR(GOES C1.1)
H2
1163.85
3.60—
193(a)
—
LW-3
(b)
Si II
1264.74 (2 lines)
4.20
149
480(b)>AR
LW-3
HI
1215.674.30
65000(b)220000(b)
>AR
LW-3
HI
1025.724.30
793(a)
6253(b)>AR
LW-2
HI
972.584.30
139(a)
1452(b)
>AR
LW-2
He I
584.344.30
580(b)
5080(b)
140000(g)LW-3 (2)Si III
1206.502
4.50
695(b)
3560(b)74000(h)
LW-3
(a)
(b)
C II
1036.340
4.55
40
290
1690(h)
LW-2
O II
718.506 (2 lines)
4.70
9.8(a)
40(a)2229(h)
LW-1
C III
977.0204.80804(a)
3666(b)58700(h)LW-2
C III
1176.0 (6 lines)
4.80
221(c)
823(d)
21310(h)LW-3
O III
703.854
4.9530(a)123(a)
4380(h)LW-1
S IV
1072.974
5.00
5.5(c)
26(a)434(h)
LW-2
N IV
765.147
5.1067(a)142(a)
5750(h)
LW-1
O IV
790.199 (2 lines)
5.15
83(b)
184(b)
8900(h)LW-1
SV786.470
5.20
27(c)
61(d)
1760(h)LW-1
NV
1238.823
5.25
61(b)549(b)
1130(h)LW-3
OV
629.732
5.40
335(b)
1019(b)90900(g)LW-3 (2)O VI
1031.914
5.50
328(a)
2460(b)8110(h)
LW-2
(b)
(b)
Ne VII465.225.75
120989
4760(h)LW-2 (2)Ne VIII
770.4285.85
54(b)600(b)
3170(h)
LW-1
(c)
(d)
Fe IX
171.073
5.90
892
8400
30600(h)SW
Fe X
174.5316.05
406(c)
5320(d)16300(h)SW
Mg X
624.962
6.05
51(b)398(b)
4190(h)LW-3 (2)
Fe XI
180.4086.15
251(c)
4380(d)13300(h)SW(c)(d)
Fe XII
195.119
6.20
174
3890
12300(h)SW
Fe XII
1241.950
6.20
3.7(c)
79(d)247(h)LW-3
Fe XIII
202.044
6.25
41(c)1248(e)3960(h)SW(c)(f)
Si XII520.666
6.25
22
1130
6790(h)LW-2 (2)Fe XIV
211.3176.30
73(c)1720(f)
10500(h)SW(e)
Ca XIV
193.9746.55—
312
766(h)SW
Ca XIV
943.587
6.55—
7.3(f)
18(h)
LW-2
Ca XV
200.972
6.65—
239(e)
751(h)SW(e)
Ca XVI
208.6046.70—
122946(h)SW
Ca XVII
192.858
6.75—
147(e)
3780(h)SW
Fe XVIII 974.860
6.80
—
88(f)
2500(h)LW-2
Fe XIX592.236
7.00—
9.3(f)
1530(h)LW-3 (2)(f)
Fe XIX1118.077.00—
8.6
1500(h)
LW-3
Fe XX
721.559
7.05——
1580(h)
LW-1
Fe XXI
786.1627.10——
178(h)LW-1
Fe XXII
845.577.10——2180(h)
LW-1
Fe XXIII 1079.414
7.15——
205(h)LW-2
Fe XXIV 192.03
7.20——
18100(h)
SW
(a)
(b)
(c)

(d)

(e)
(f)

(g)
(h)

From SUMER spectra. The active region (AR) values are from a modest one.
Vernazza and Reeves, 1978, ApJS 37, 485.
CHIANTI 7.0 (CHIANTI ionization equilibria and coronal abundances, density=5×109 cm −3, quiet Sun
DEM).
CHIANTI 7.0 (CHIANTI ionization equilibria and coronal abundances, density=5×109 cm −3, active region
DEM).
Warren et al., 2011, ApJ 734, 90.
CHIANTI 7.0 (CHIANTI ionization equilibria and coronal abundances, density=5×109 cm −3, active region
core DEM from Warren et al. 2011).
From CDS spectra of a GOES C1.1 flare.
CHIANTI 7.0 (CHIANTI ionization equilibria and coronal abundances, density=5×109 cm −3, DEM from CDS
spectra of a GOES C1.1 flare).

29/1006.2.3 Radiometric performances to achieve high-throughput optics
The instrument must achieve high throughput performance. The effective area (EA) in cm2 is related
to the radiance, L, in erg/scm2sr by the following formula:𝑁[cm2 ] ,𝐸𝐴 = 8.3 × 102
𝐿𝑡𝑒𝑥𝑝 𝜆where N is photon number, texp is exposure time in second, and 𝜆 is wavelength in Å. EA as a functionof wavelength was calculated and is given in Figure 6.2(a) for SW and LW bands based on the EAsof Hinode/EIS and SoHO/SUMER, respectively. The EA of the SW band is calculated under theassumption that the diameter of the primary mirror is twice as large as Hinode/EIS and the visible-IRfilter in the optical path is removed. The EAs of LW bands are calculated under the assumption thatthe number of reflections is reduced from 3 to 1 and the diameter of the primary mirror increases
from 12~cm to 28~cm. The calculation includes the mirror, grating and detector efficiencies, as wellas the relevant geometrical factors (i.e., mirror area and the splitting of the grating into the SW and
LW channels). The plot includes, for comparison, the published effective areas of Hinode/EIS,SoHO/SUMER and SoHO/CDS. The peak efficiencies give a factor of ≈ 10 improvement withrespect to Hinode/EIS in the EUV and an improvement of a factor of ≈ 40 over SoHO/SUMER in the
FUV. Figure 6.2(b) shows the expected signal (in count/arcsec2) obtained by multiplying the EA by
the radiances for different solar regions and by considering typical exposure times for those regions
and a resolution element of 1~arcsec2.
(a)

(b)

Figure 6.2: (a) Assumed Solar-C_EUVST effective area based on the baseline architectures described in section 15.
The effective areas of Hinode/EIS, SoHO/SUMER, and SoHO/CDS are also shown for comparison. (b) Expected countrates (count/arcsec2) for the indicated exposure times for different solar observational targets (5~s for the quiet Sun, 1~sfor active regions, and 0.5~s for a small flare). The horizontal dashed line marks the 200~counts (in the spectral line)
level necessary to determine line positions with a ≤ 2~km/s accuracy as shown in section 6.2.4.

6.2.4 Photometric accuracySpectroscopic observations are challenging because they distribute the observed emission into manysmall spectral bins. Furthermore, spectroscopic plasma diagnostics often rely on high-order momentsof the line profile or on information from multiple lines. To evaluate the diagnostics of interest we
have performed Monte Carlo simulations where we assume a total line intensity, Doppler shift, nonthermal velocity, and background and generate a random realization of the line profile. We then infer
the values of these parameters from a least-squares fit to the synthetic data, just as we will for theactual observations. Repeating this process for different intensity levels allows us to estimate how theuncertainties in the diagnostics depend on the observed counts.
An example calculation is shown in Figure 6.3. Here we have simulated the C~III 977.02~Å line
assuming no Doppler shift, a non-thermal velocity of 30~km/s, a dispersion of 37~mÅ per spectral
pixel, an instrumental broadening (full width at half maximum; FWHM) of 2.5~pixels, and a
30/100

background level of 2.5 counts per
spectral bin. This simulation
indicates, as one would expect, thatapproximately 100 total counts in theline profile are needed to measure theintensity to 10%. An error of 2 km/s inthe Doppler shift is achieved forapproximately 200 total counts.
Approximately 150 counts are neededto achieve an uncertainty of 15% in
the non-thermal velocity. Finally, we
have also estimated the uncertainty inthe electron density, which is derived
from the ratio of two lines, as afunction of total counts. Here we use
the Fe XII 186.867/195.119 Å ratio to
illustrate this. About 700 counts are Figure 6.3: Errors in the moments of a line profile as a function of total
counts. For the electron density the counts are those for the weakerneeded to measure the electron line.
density to 30%. Note that the needed
count levels depend strongly on the line width. Lines formed at higher temperatures will have a larger
thermal component in the width and require more counts to achieve the same level of accuracy in the
Doppler shift and non-thermal velocity.
~

~

~

~

Another source of uncertainty in the plasma diagnostics is the instrument calibration. Of primaryconcern to Solar-C_EUVST is the relative calibration, which impacts the use of diagnostics formed
from emission observed at different wavelengths. Solar-C_EUVST will have a relative calibration of
25% over all wavelengths. This would increase the uncertainty in the density to a factor of 4 (0.6 inthe log) for lines widely separated in wavelength. However, for lines close in wavelength the relative
uncertainty in the calibration will generally be much less and the uncertainty will be lower.
6.2.5 Comparison with other UV spectroscopic instruments
Observation performances of current and forthcoming coronal UV spectrometers introduced in
Section 3.3 are compared in Table 6.3. The enormous number of spectral lines with high spatial andtemporal resolutions recorded by Solar-C_EUVST will allow determination of the thermodynamic
state of the plasma (density, temperature) and its chemical composition at all temperatures.
Observations of this completeness are unprecedented.

6.3.

Major flight / ground trades [6-3-3]")53

5

Table 6.4 gives the results from our conceptual study. The study suggests that the following items
shall be studied in more details in the design phase:• The best method to capture large flares with a rather narrow field of view instrument is semireal-time pointing adjustment operations. The baseline method will be to follow that used in
Hinode’s flare watch operations which allows for such last-minute pointing adjustment.
• The amount of data will increase during flare campaigns because of the inherently unpredictable
nature of flare onset. Appropriate methods for storing and downloading flare data volumes will
be studied.
31/100Table 6.3: Solar-C_EUVST performance in comparison with current and forthcoming spectrometers. Yellow filledcells show the specifications that meet our observation target.
Existing mission
Mission Under Study
Hinode/SoHO/
Solar Orbiter/
Solar-C_
IRIS
MUSE
EIS
SUMER
SPICE
EUVST
10 – 60 s
<10 s
1.0 s
0.2(4) – 5 s
Temporalfor AR
for AR
for AR
for AR
resolution for a≥15 – 60 s (2)5s
20 – 120 s
10 s
0.1 s
0.2 – 20 s
(1)slit position
for QSfor QS
for flare
for QS
Spatial resolution(1″=725 km on
the Sun)2″ – 3″

1.5″

0.33″

1″

0.4″

0.4″

Field of View
(without
repointing)±290″×512″±1920″×300″

32″×120″960″x240″

170″x170″±150″×280″

(1)
(2)

(3)
(4)

6.4.

Wavelength
bands [Å]170 – 210
250 – 290

500 – 1600

1332 – 1358
1389 – 1407
2783 – 2835

702 – 792
972 – 1050
485 – 525

107 – 109
170 – 174276 – 288

170 – 215
463 – 542
557 – 637
690 – 850
925 – 1085
1115 – 1275

Primarytemperature
coverage(log10 T/[K])4.9, 5.6 – 7.2

3.6 – 5.8,
6.8 – 7.0(3)3.7 – 5.2,
7.0

3.6 – 5.8,
6.8 – 7.0

5.9, 6.4,
7.0-7.1

4.2 – 7.2

Diagnostics
capabilities

LOS Velocity
LOS Velocity
ElectronElectronDensityDensity
Electron/ion
Electron
Temperature
Temperature
DEM

LOS
velocity

LOS Velocity
ElectronDensity
Electron/ion
Temperature

LOS
velocity

LOS Velocity
ElectronDensity
Electron/ion
Temperature
DEM

The values given here are typical exposure duration for multiple numbers of spectral lines covering a wide
temperature range.
There are several lines in the SoHO/SUMER range that could actually be recorded at higher cadence even on
the quiet Sun but there is neither sufficient telemetry nor on-board storage to allow this. Values given here
refer to the time required to transmit just one window of 50×300 pixels (300″slit). None of these problems
will affect Solar-C_EUVST.
SoHO/SUMER can record spectral lines in only one spectral window of ≈40 Å at a time. If spectral lines inanother spectral window are measured, the spectral window must be changed before the exposure.
Exposure time and time step of 0.2 s is for one line per camera (4 spectral lines).

Data to be returned in the course of the investigation (= mission data): telemetry data, sample
data etc. [2-2-1]
ミッションで獲得するデータ, テレメトリーデータ，サンプルなど (以下ミッション
データ)

This section describes the data to be returned from the proposed mission instrument. Typical examplesof observing sequences required for achieving each science objective are shown in section 6.4.1 todefine the observing modes and estimate the amount of data, followed by the requirement for
telemetry data in section 6.4.2 and the expected data analysis in section 6.4.3.

32/100Item
Decision of theobservation target
and the timeline

Handling the large
amount of the datafrom observations

Onboard processes

Observations oftransient phenomena

Estimation of theorbital variation of
the wavelength shifts
Calibration of
wavelength drift due
to the temperature
variations
Adjustment of the
focus position

Radiometriccalibration

6.4.1

Table 6.4: Flight and ground trades
Results from conceptual study• Decided and confirmed by observers in the daily meeting. The observers monitor
solar activities by using full disk images provided by other ground-based/space-bonetelescopes. Once the targets are confirmed, observers make a timeline and decide
observation sequences for each target. The decided observation plans are uploaded
to the satellite. This is what has been done in Hinode daily operations.
• The data downlink capability may be insufficient to download the larger amount of
the data to the ground.
• Depending on scientific target, observers choose the detector areas that include theappropriate bright lines and the size of the field of view during the planning of theobservations. Thus, only parts of the detector areas are downloaded.
• Lossy and lossless compressions of the observed data are performed in the onboard
processor• Higher order data products (e.g. intensity, velocity, line width) will not be derived
onboard. This is the same as Hinode.
• To increase chances to observe transient phenomena, semi-real-time pointing
adjustment operations or large data memory should be installed.
• The semi-real-time pointing adjustment will be performed manually by the ground
operation.
• An onboard function may be required to protect the science data when large flares
occur.
• To confirm preserved data to be downloaded, thumbnail images are provided to theground operation.
• By deriving relative wavelength among lines emitted from the same ion species, theshift can be calibrated in the data analysis on the ground. The wavelength shift istypically ± 6 km/s in the case of a sun synchronous orbit.
• An artificial neural network is incorporated to establish a relationship between the
instrumental temperatures and the spectral drift. The same method has been used for
the Hinode/EIS wavelength calibration (Kamio et al. 2010).
• The best focus position is evaluated in the initial phase of the flight by analyzing
data taken at several focus positions on the ground.
• Focus mechanisms are required to be onboard based on the tolerances required forthe optics optimum performance.
• The opt-mechanical design should be taken so that the focus change is negligible forthe orbital variation.
• Radiometric calibration is performed by using an internationally recognized sourceon the ground.

Required observations

One way to think about the design requirements for the telemetry data is to devise a set of basic
observing modes that can be used singularly or in combinations (observing sequences or programs)
to achieve the instrument and mission science objectives. Table 6.5 provides typical examples of the
programs (observing sequences) expected to be used for each of the science objectives. The table also
gives an estimate of the duration of one observation required for each study, and the data volume
produced by each study when the study is run in the duration.
The basic observing modes used to form the studies described in Table 6.5 are listed in Table 6.6.
Collectively, they define the design requirements for the instrument. The estimated Solar-C_EUVST
data rate production of each observing program and mode is given in Tables 6.5 and 6.6, respectively.

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Science
Tasks ofsub-objective observations
I-1:
I-1-1:
NanoflaresI-1-2:
I-1-3:

I-1-4:
I-2:
Wave

I-2-1:
I-2-2:
I-2-3:

I-3:
SpiculesI-3-1:
I-3-2:

I-4:
Solar Wind

I-4-1:
I-4-2:

II-1:
II-1-1:
Fast Magnetic
Reconnection

II-1-2:II-1-3:
II-2:
Flare trigger

II-2-1:II-2-2:

Table 6.5: Required observing sequences (programs).
Observing sequences
Duration / Data volume
1) AR-Context × 1
87.5 s / 0.11 Gb
2) AR-High_S&S × 100
300 s / 0.17 Gb
Repeat 1) & 2) × 9
3487.5 s / 2.52 Gb
1) AR-Context × 1
87.5 s / 0.11 Gb
2) AR-High_S&S × 100
300 s / 0.17 Gb
Repeat 1) & 2) × 9
3487.5 s / 2.52 Gb
1) AR-Context × 1
87.5 s / 0.11 Gb
2) AR-High_Small × 1
250 s / 0.06 Gb
3) AR-Mid_Small × 14
525 s / 0.31 Gb
Repeat 2) & 3) × 4
3187.5 s / 1.92 Gb
1) AR-Context × 1
87.5 s / 0.11 Gb
2) AR-Mid_Small × 14
525 s / 0.31 Gb
Repeat 1) & 2) × 6
3675 s / 2.52 Gb
1) AR-Context × 40
3500 s / 4.12 Gb
2) AR-Mid_Small × 96
3600 s / 2.12 Gb
1) or 2) × 1
3500 (3600) s / 4.12 (2.12) Gb
1) AR-Context × 1
87.5 s / 0.11 Gb
2) AR-High_S&S × 100
300 s / 0.17 Gb
Repeat 1) & 2) × 9
3487.5 s / 2.52 Gb
1) AR-Context × 1
87.5 s / 0.11 Gb
2) AR-High_Small × 41000 s / 0.25 Gb
Repeat 1) & 2) × 5
5437.5 s / 1.8 Gb
1) QS-Context × 1
262.5 s / 0.21 Gb
2) AR-High_SSmall × 25
1250 s / 0.88 Gb
Repeat 1) & 2) × 2
3025 s / 2.18 Gb
1) QS-Context × 1
262.5 s / 0.21 Gb
2) AR-High_SSmall × 25
1250 s / 0.88 Gb
Repeat 1) & 2) × 2
3025 s / 2.18 Gb
1) AR-Context × 1
87.5 s / 0.11 Gb
2) AR-Mid_Full × 2
1050 s / 0.62 Gb
Repeat 1) & 2) × 3
3412.5 s / 2.19 Gb
1) QS-Context × 1
262.5 s / 0.21 Gb
2) QS-Low_S&S_L × 1
300 s / 0.01 Gb
3) QS-Low_Small_L × 41200 s / 0.05 Gb
Repeat 1) & 2) & 3) × 2
3520 s / 0.54 Gb
1) AR-Context × 1
87.5 s / 0.11 Gb
2) FL-Mid_Full × 4
700 s / 1.23 Gb
3) FL-Mid_Small × 56
700 s / 1.23 Gb
4) FL-Mid_S&S × 1400
700 s / 0.62 Gb
Repeat 1) & 2) × 240
2 days / 322 Gb
Repeat 1) & 3) × 240
2 days / 322 Gb
Repeat 1) & 4) × 240
2 days / 175 Gb
1) AR-Context × 1
87.5 s / 0.11 Gb
2) FL-Mid_Small × 56
700 s / 1.23 Gb
3) FL-Mid_S&S × 1400
700 s / 0.62 Gb
Repeat 1) & 2) × 240
2 days / 322 Gb
Repeat 1) & 3) × 240
2 days / 175 Gb
1) QS-Context × 1
262.5 s / 0.21 Gb
2) AR-Mid_Small × 14
525 s / 0.31 Gb
Repeat 1) & 2) × 43150 s / 2.08 Gb
1) AR-Context × 1
87.5 s / 0.11 Gb
Repeat 1) every 30 min 336 times 1 week / 36.96 Gb
1) AR-Context × 1
87.5 s / 0.11 Gb
2) AR-Mid_Small × 30
1125 s / 0.66 Gb
Repeat 1) & 2) × 144
2 days / 111 Gb

34/100

Data rate[Mbps]
0.72
0.72

0.60

0.69
1.18
(0.59)
0.72
0.33
0.72
0.72
0.64

0.15

1.87
1.87
1.01

1.87
1.01

0.66
0.06
0.64

Table 6.6. Observing modes
Obs. mode
FOV
Slit width /Step Cadenc
T rangeN°
Data rate(1)exposure
size
e
lines
1
QS-Context
280″×280″
1.6″ (1.5 s)
1.6″ 262.5 s 0.03 – 1.5 MK
16
0.78 Mbps
2
QS-High_Full
100″×100″ 0.40″ (15 s) 0.40″ 3750 s 0.03 – 1.5 MK
8
0.06 Mbps
3
QS-High_Small
20″×100″
0.40″ (15 s) 0.40″
750 s 0.03 – 1.5 MK
8
0.06 Mbps
4
QS-Mid_Full
100″×100″
0.80″ (5 s)0.80″ 1750 s 0.03 – 1.5 MK
8
0.24 Mbps
5
QS-Mid_Small
20″×280″0.80″ (5 s)
0.80″
125 s 0.03 – 1.5 MK
8
0.24 Mbps
6
QS-High_S&S
0.4″×100″0.40″ (9 s)0.40″
9s0.03 – 15 MK
8
0.09 Mbps
7
QS-Mid_S&S
0.8″×140″
0.80″ (5 s)0.80″
5s0.03 – 15 MK
8
0.12 Mbps
8
QS-Low_S&S_L
0.8″×280″0.8″ (300 s)
0.8″
300 s 0.03 – 15 MK
160.01 Mbps
9
QS-Low_Small_L
16″×280″
1.6″ (30 s)
1.6″
300 s 0.03 – 15 MK
160.04 Mbps
10 AR-Context*
280″×280″
1.6″ (0.5 s)
1.6″
87.5 s 0.03 – 15 MK
16
1.18 Mbps
11 AR-High_Full
100″×100″0.40″ (5 s)0.40″ 1250 s 0.03 – 15 MK
120.25 Mbps
12 AR-High_Small
20″×100″
0.40″ (5 s)
0.40″
250 s 0.03 – 15 MK
120.25 Mbps
13 AR-High_SSmall
10″×100″
0.40″ (2 s)
0.40″
50 s
0.03 – 15 MK
120.63 Mbps
14 AR-Mid_Full*280″×280″ 0.80″ (1.5 s) 0.80″
525 s 0.03 – 15 MK
120.59 Mbps
15 AR-Mid_Small*
20″×280″0.80″ (1.5 s) 0.80″ 37.5 s 0.03 – 15 MK
120.59 Mbps
16 AR-High_S&S
0.4″×100″0.40″ (3 s)0.40″
3s0.03 – 15 MK
160.56 Mbps
17 AR-Mid_S&S
0.8″×140″ 0.80″ (1.5 s) 0.80″
1.5 s
0.03 – 15 MK
160.78 Mbps
18 FL-Mid_Full*280″×280″ 0.80″ (0.5 s) 0.80″
175 s 0.03 – 15 MK
121.76 Mbps
19 FL-Mid_Small*
20″×280″0.80″ (0.5 s) 0.80″ 12.5 s 0.03 – 15 MK
121.76 Mbps
20 FL-Mid_S&S*
0.8″×140″ 0.80″ (0.5 s) 0.80″
0.5 s
0.03 – 15 MK
120.88 Mbps(2)
21 Slit camera monitor
30″×150″(0.2 s)n/a7s
0.01 MKn/a
0.06 Mbps(3)(1)
Based on 60 pixel wide windows. The number of pixels along the slit is given by the y-size of the FOV divided
by the slit width (spatial binning equal to resolution across slit). 14 bits/pixel and lossless compression (factor
2) are assumed. Note that lossless compression of EIS is ~35%, largely depending on the features in the data.
We assume lossy compression (factor 4) in the case of flare and some AR study, marked *.
(2)
Solar-C_EUVST context and Slit camera monitor images are also needed to facilitate alignment with otherobservations (e.g., DKIST).
(3)
Compression to 2.2 bits/pixel is assumed.

6.4.2 Telemetry and onboard data recorder requirements
Solar-C_EUVST’s broad temperature coverage (multiple lines to be recorded simultaneously) at hightemporal resolution requires an adequate telemetry rate to ensure operations. According to Table 6.5,
the data rates expected from each observing program of most tasks are up to 0.72 Mbps. The typicalduration of each observing program in Table 6.5 is about 1 hour, and a variety of observing programs
run during the daily observations to optimize the science return as described in section 11. Thus, whenobservations are carried out continuously without idle periods, an average data rate during theobservations is about 0.72 Mbps in the most programs. This is the minimum requirement for the data
rate. If a lower data rate is required due to restrictions on the telemetry resource in some periods, idle
times will be imposed between two consecutive observing programs.
~~

A data recorder is required to store the science data from the instrument, because continuous contacts
with ground stations for 24 hours are unlikely (trade-off study in section 14.2 (6)). Moreover, Tasks
I-2-1, II-1-1, and II-1-2 require higher data rates of 1.18, 1.87, and 1.87~Mbps, respectively. Science
data obtained at these data rates should be stored adequately in the memory of the onboard data
recorder. For Task I-2-1, a volume for storing the produced data, 4.12 Gbits (1.18 Mbps × 1 hour), is
required in the data recorder. Regarding the flare-related observations of Tasks II-1-1 and II-1-2, it isdesirable to continuously run the defined observing programs as long as possible since the predictionof flare occurrence is difficult. The continuous observations will produce 162 Gbits in a day (1.87Mbps × 24 hours). However, data that are truly necessary for achieving Tasks II-1-1 and II-1-2 are
obtained typically within 3 hours around the onset of the flare, and the amount of the data is 20 Gbits(1.87 Mbpts × 3 hours). Therefore, 20 Gbits is the minimum requirement for the data recorder volume.
~

~

~

~

~

35/100Note that, as listed in section 6.3, an onboard function to protect flare data from being overwritten ora quick-look browser before data downloading would be required.
In summary, the minimum requirements of the data rate and the volume of the onboard data recorder
are 0.72~Mbps and 20~Gbits, respectively.
6.4.3 Usage of the data
From the spectroscopic data obtained by Solar-C_EUVST, various plasma parameters such as density,
temperature, and elemental abundances are derived by reference to atomic data base like CHIANTI(Dere et al. 1997; Del Zanna et al. 2015). Figure 6.4(a) shows the density range covered by suitabledensity-sensitive line ratios available to Solar-C_EUVST as a function of the line formation
temperature, whereas Figure 6.4(b) shows the temperature range covered by suitable temperaturesensitive line ratios available to Solar-C_EUVST as a function of the line formation temperature.
Non-thermal velocities are determined by subtracting the thermal and instrumental widths from theobserved line width. Doppler velocities in the long wavelength channels are generally measured
relative to emission lines from neutral atoms, which are assumed to be at rest. In the short channel,
Doppler velocities are measured relative to the quiet Sun and have larger uncertainties.
Comparisons with simultaneous photospheric and low-chromospheric measurements obtained byother space-borne and ground-based observatories (e.g., magnetograms by SDO/HMI and DKIST)
are essential for revealing the relationship between the atmospheric dynamics and surface evolutions.
Furthermore, various kinds of state-of-the-art numerical simulations support the interpretation ofobservational results. The proposed scientific objectives will be achieved through the balanced
combination of the different methods mentioned above.
Data analysis tools, including the software for data prepping and fitting Gaussian line profiles, will
be provided by the Solar-C_EUVST science team. Observational data will be publicly available soonafter initial calibrations are made. Online search tools that help users locate data will also be provided.

(a)

(b)

Figure 6.4: (a) Electron density range covered by the density-sensitive intensity line ratios available in the
Solar-C_EUVST spectral bands as a function of the formation temperature of the line pair. Densities
obtained from ratios with all components on the same channel are shown in blue while those having
components on different channels are shown in red. (b) Electron temperature range covered by thetemperature-sensitive intensity line ratios available in the Solar-C_EUVST spectral bands as a function of
the formation temperature of the line pair. The best diagnostics (bright or relatively bright lines and goodtemperature coverage) are shown in blue. Density and temperature ranges are evaluated with CHIANTI 7.1.

36/100

7.

Scientific traceability matrix (draft) [1-4-1]
( [eiT\cZO_[cYQW

The science requirements traceability matrix (Table 7.1) shows the breakdown flow from scienceobjectives to key instrument and mission data requirements. The driving requirement forinvestigations (third column from the left) is based on Tables 5.1 to 5.6. If one sub-objective
proposes multiple observing tasks, we show the most demanding values for the driving requirement.
These driving requirements are then translated into a set of instrument requirements (second column
from the right) in the following manner. Broad temperature coverage in bright emission lines
directly determines the wavelength requirement. The effective area is derived from the required
spatial resolution, temporal resolution, and count rates; the count rates here are determined so as tomeet the velocity (Doppler and width) resolution requirements. These calculations include a
combined spectral width of 2.5 pixels, which considers the slit width and instrumental broadening(see section 6.2.4). However, the sensitivity study using a combined spectral width of 3.0 pixels
show only a minimal (~20%) increase in the required amount of counts. Therefore, 2.0 pixels wereadopted as the spectral resolution requirement. We defined the effective area for each wavelengthband (SW, LW1, LW2, LW3) as the area necessary to obtain 200 counts per exposure. Here, thevalue of 200 comes from the requested velocity resolution (e.g., 2 km/s) in the driving requirement
for investigations (third column from the left), while the exposure time is based on the temporal
resolution in the same column. We calculated the count rates for the typical emission line of eachwavelength band, i.e., Fe XII (SW: 195 Å), Ne VIII (LW1: 770 Å), Fe XVIII (LW2: 975 Å), and C
III (LW3: 1176 Å). Observing program requirements determine the FOV- and raster-related
requirements. Data and compression requirements as well as downlink capabilities are directly
driven by the characteristics of the observing program. Co-registration of the slit position is needed
to analyze the observations in conjunction with other data sets.

37/100Table 7.1: Scientific traceability matrix

38/100Table 7.1: Scientific traceability matrix (continue)

1) 1": 730 km, Vd: Doppler velocity, Vnt: nonthermal velocity (e.g., turbulence)2) Coregistration of the slit position with solar features of <0.18”.
3) Required spectral resolution corresponds to a two pixel average FWHM over the central 100” of eachdetector.
4) Required spatial resolution is given in FWHM.

39/100