Chapter 2:

Solar Ultra-violet Visible and IR Telescope

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

2-1

2.1.2. Science targets with the instrument


Magnetized solar atmospheres: how they are created and evolve.

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

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

Small-scale dynamic phenomena responsible for the energy transfer.

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

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

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

Polarimetric sensitivity and accuracy

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

2-4

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

Spatial resolution

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

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

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

Field-of-view

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

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

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

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

1.2λ/D[arcsec]

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

5

Response height
[km]
Te
V

Zeeman effectHanle
effectgeff

Blong [G]Btrans [G]

Mg II k 2796 Å
Mg II h 2803 Å

0.046
0.046

2.5x102.8x1042.6x102.2x105

~1300
~1200

~1400
~1300

1.17
1.33

85
73

2100
1800

X
-

Ca II K 3933 Å
Ca II H 3968 Å

0.065
0.065

5.5x1056.4x1052.9x105
3.0x105

~1000
~900

~1200
~1100

1.17
1.33

62
51

1500
1200

X
-

H I β 4861 Å
Mg I b3 5167 Å

0.080
0.085

1.1x107
1.3x1071.8x106
1.3x106

~1500
~500

~1000
~600

1.00
2.00

16
3.7

1200
220

X
X

Mg I b2 5172 Å
Mg I b1 5183 Å

0.085
0.086

1.3x107
1.3x1071.0x106
9.2x105

~500
~500

~600
~600

1.75
1.25

6.5
9.6

250
360

X
X

Na I D2 5890 Å
Na I D1 5896 Å

0.097
0.097

1.5x1071.5x1077.2x1058.4x105

~400
~400

~600
~500

1.17
1.33

4.7
3.4

290
230

X
-

H I α 6563 Å
Ca II 8498 Å

0.108
0.140

1.5x1071.1x1072.8x106
3.3x106

~1500
~700

~1300
~1000

1.00
1.07

9.7
6.3

680
300

X
X

Ca II 8542 Å
Ca II 8662 Å

0.141
0.143

1.1x1071.1x1072.1x106
2.1x106

~800
~700

~1200
~1100

1.10
0.83

6.7
8.4

340
430

X
X

He I 10830 Å (QS)He I 10830 Å (AR)

0.179
0.179

1.8x1071.8x1071.7x1071.5x107

~1800
~1800

~1800
~1800

1.42
1.42

18
9.3

380
270

X
X

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

2-7

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

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

2-8

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

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

Blong

0.1 arcsec sampling
1 sec integration~20 gauss

Btrans~380 gauss (Zeeman)Blong~10 gauss

Btrans~270 gauss (Zeeman)

Blong
Btrans~7 gauss~340 gauss (Zeeman)

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

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

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

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

2)

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

2-10

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

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

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

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

SP
x

BF
x
x
x
x

x
x

NF

x
x

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

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

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

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

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

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

2-13

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

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

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

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

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

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

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

Ring plate

Center section

Top ring
Upper truss

Lower truss

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

2-16

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

Despace

M2

Decenter
Tilt
Despace
Decenter
Tilt

1)

2)
3)
4)

Aberration
DefocusSpherical
Coma
Coma
DefocusSpherical
Coma
Coma

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

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

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

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

2-17

M1

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

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

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

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

Absorption coeff.
α
0.12
0.1
0.12
0.11
---

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

2-19

Absorbed radiation
(W)
244.8
175.4
9.0
2.1
431.4

300

1940

Φ2233
shield tube
M1-M2
2800

HDM
cylinder

Opticalbench unit

2931
2296

Φ60
truss
pipe

1680
1680

Φ2120
Lower tube

300

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

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

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

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

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

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

2-20

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

Absorption coeff.
α
0.17
0.1
0.17
0.11
---

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

Absorbed radiation
(W)
348.6
165.4
12.1
2.0
528.2

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

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

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

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

2-21

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

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

2-22

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

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

854 nm
Ca II
19
34 mÅ

280 nm
Mg II
58
11 mÅ

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

2-23

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

Normal mode

Deep mode

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

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

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

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

Pass-bands
Spectrum lines

279 nmMg II k/h
2-24

380 nm
CN band

393 nmCa II K

TBD
Contin.

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

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

2-25

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

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

Pass-bands
Spectrum lines

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

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

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

2-26

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

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

2-27

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

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

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

2-28

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

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

interferometer

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

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

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

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

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

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

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

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

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

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

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

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

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

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

- a5 -

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

- a6 -