�IRIS ＋ ひのでInterface Region Imaging Spectrograph‣ NASA 小型衛星ミッション (SMEX)‣ 2013年6月27日打ち上げ‣ ひのでの観測を踏まえて提案・開発‣ 紫外線分光観測 → 物理量の取得（高空間分解能）● 2次元面運動＋● 視線方向速度● 非熱的速度● 多温度観測

彩層の役割を理解する

�IRIS 観測波長と観測例

we get spectral lines in emission and absorption?

太陽表面
 ⇥



 

 ⇤









Mg II 吸収線











源泉関数


Mg II





Han Uitenbroek/NSO

スリット
方向





I Basic Radiative Transfer

スリットジョー動画（Mg II k）

スリット位置Mg II k

波長

Mg II h�IRIS 観測波長と観測例

Liu+2014 in prep.

フレアと噴出スリットジョー（C II）
スリットジョー（C II）

Mg II (2796A)
10,000 K

Mg II kMg II hC II (1336A)
20,000 K

Si IV (1403A)
80,000 K

�Science IRIS 特集号
B. De Pontieu→ 馬場

On the prevalence of small-scale twist in thesolar chromosphere and transition region

H. Peter→ 勝川

Hot explosions in the cool atmosphere if theSun

V. Hansteen
The unresolved fine structure resolved: IRIS
observations of the solar transition region

P. Testa

5本の論文（10月17日号）

Evidence of nonthermal particles in coronalloops heated impulsively by nanoflares

H. TianPrevalence of small-scale jets from the
networks of the solar transition region andchromosphere

T

Viewing the same limb with the Solar Dynamspectra, with models of the energy and mass
ics Observatory Atmospheric Imaging Assembly
exchange between the cooler chromosphere and(SDO/AIA) instrument (12) in the 304 Å channel
hot corona. For one, the observed intensities of(which is dominated by He II, at 100,000 K) andlower transition-region lines are much greater
the coronal 171 Å (Fe IX/X at ~106) and 193 Å (Fethan can be accounted for by thermal conducXII at ~1.5 106 K) channels (13), we do not clearlytive flux flowing back from the corona. Furobserve UFS-related structures (Fig. 1 and movie
thermore, lower transition-region lines show,
S3). This is for two reasons: (i) the AIA spatial
on average, Doppler redshifts on the order of1
Institute of Theoretical Astrophysics, University of Oslo,
resolution is insufficient to resolve the structures
10 km/s (1), one-third the speed of sound. Based
Post Office Box 1029, Blindern, NO-0315, Oslo, Norway.
discussed here, and (ii) bound-free absorptiononindirect
spectroscopic
evidence
fromHigh2
E YE IN G TH E S U N
Lockheed Martin Solar and Astrophysics Laboratory, 3251
by neutral hydrogen and singly ionized helium
Resolution Telescope Spectrograph (HRTS) andHanover Street, Org. A021S, Building 252, Palo Alto, CA
3renders the AIA opaque to extreme ultraviolet
Skylabspectra,itwasalready
postulatedin1983
94304, USA. Harvard-Smithsonian Center for Astrophysics,
60 Garden Street, Cambridge, MA 02138, USA. 4Department
therebyasshielding
theoflower
thedown
dominant
emissionfrom posurelines toformedchangesthat
on scalesto the shortest
cadence
exposure, canemission,
remain recognizableaa cadence
54 s andtransition
is an example of a fairly longof Physics, Montana State University, Bozeman, Post Officedata inspectedsofar (everyregion4 s), examples
ofsystem overmagperiods extendingtoseveraltens ofthelived“nest”
of loops
that remains active during
regionandmakingUFS
nearlyinvisible.
inthetransition
occursinstructures
Box 173840, Bozeman, MT 59717, USA. 5Bay Area
which are shown in region of interest (ROI) 2 andminutes. The temporal evolution
of
the
three
rethe
entire40-minspanthatloops
the observations lasted.
One of the most striking features of UFS
netically isolated from the corona called theEnvironmental Research Institute, 596 1st Street West, ROI 3 of Fig. 1. On the other hand, a system ofgions of interest are shown in Fig. 1 (more loops
Movies S1, S2, and S3 show these nests to be comtemporalvariability.
Isolated
UFSloops
loops
fineloops
structure”(UFS)
However,
Sonoma, CA 95476, USA.
loops, in“unresolved
which individual
vary fromex- (2–5).
are displayedin fig. S1).isROI
1 was observed
with
posed
of many
withlight
more or less cospatial*Corresponding author. E-mail: viggoh@astro.uio.no
up, either partially or footpoints
wholly, that
andlight
show
large
the following decades have not brought conup anddarken episodically.

he outer solar atmosphere between the
104 K chromosphere and the 106 K corona,
the so-called transition region, has long
puzzled solar physicists (1). It has been difficult to reconcile measured intensities andmotions, either directly observed or inferred from

�Science IRIS 特集号

Hansteen+2014

いままで空間分解されていなかった微小ループ構造が初めて分解された→ 短いループの急激な加熱・冷却が効きすぎるためコロナ温度に達しない・
����3次元モデルで再現できていない高速成分（70 km/s）がある
Fully formed loops are seldom seen. Rather, loops
appear to be lit up in segments, with each segvisible for roughly a minute. There
23:51:19
23:53:12
23:55:05 ment only being
23:56:57
0.20
is also a tendency for the UFS loops to appear to
0.15
rise with time, as can be seen in ROI 2 (Fig. 1).
We find that the UFS loops have a full length
0.10
1
of 4 to 12 Mm (106 meters), a maximum height2
0.05
of 1 to 4.5 Mm with an average of 2.5 Mm, and amedian intensity in data numbers per second of
0.00
40 to 50 DN/s (Figs. 2) (14). This intensity is
0
4
6
8
10
12
2
4
6
8
10
12
14
16
32
23:52:53
23:53:12
23:53:31 larger than 23:53:49
the measured intensity of the backApparent Length [Mm]
Projected Full Length [Mm]ground “spicules”—longer nearly radial features—
0.25
of 15 DN/s, which is also apparent from visualinspection of Fig. 1. Although the detailed fill0.20
ing factor of either component is not well known
0.15(given the superposition at the limb and theirlimited visibility on the disk), it is clear that both
0.10
23:52:53
23:53:12
23:53:31 resolved components
23:53:49 have an important role in
0.05
the lower transition-region emission, with the relative contribution dependent on the local mag0.00
0
2
4
6 0
20
40
60
80
100
120 netic field topology.
By aligning the slit along the limb, IRIS also
Maximum Height above Limb [Mm]Median Intensity [DN/s]
allowsone to gather spectral data of the UFS
Fig. 2. UFS loops are short, bright, and low-lying. Properties of 85 UFS loops taken from a limb
(example in Fig. 3). The spectrum shows large
observation data set: (A) apparent length of the visible loop segments, (B) projected full length along
excursions as a function of position along theloop, (C) maximum height, and (D) median intensity of the visible loop segments. The intensity
loop, implying large plasma velocities, toward
measured for the dimmer “spicular” background is shown with a dashed line.
the red as well as toward the blue. We find extremeline loops
profilesonat theloop footpoint—theFig. 1. IRIS Si IV 1400 Å slit jaw images reveal highly dynamic, low-lying loops at transition-region temperatures. (A) UFS
the upper
western
portion
of
the
loop
that
meets
solar limb. The slit is evident as a dark line near solar-x = 971 arcsec. (B) The same field of view is shown, but with SDO/AIA images: the coronal 171 Å the underlying
atmosphere—during the entire 200-s lifetime of(blue) and 193 Å (green) filters (movie S3 shows12the same field of view, but with the AIA He II 304 Å filter). The rapid evolution
ofthe
theUFSUFS
in withthree
ROIs excursions
is
12
loop,
redwardof 70 todemonstrated in the small panels on the right sideofthefigure.
The
time
of
each
exposure(hour:minute:second)isindicated
at
the
top
of
the
panels.
80
km/s.
This
is
two
to
three
times the speed of1010sound in a 80,000 to 100,000 K plasma. We
88observed the spectral properties of several UFS
loops and find 1255757-1
that such high velocities occur
SCIENCE sciencemag.org
17 OCTOBER 2014 • VOL
6
6 346 ISSUE 6207
often, although not always. This indicates that
the UFS loops are locations of episodic and44violent heating.
22
How can we put these observations into thecontext of coronal heating models and the struc00ture of the upper solar atmosphere? Models that
0
5
10
15
−150 −100
−50
0
50
100
150
arcseckm/s
assume most transition-region emission stems
Afrom loops connected to the corona, and there1.0
100
fore whose temperature structure is maintainedby thermal conduction, lead to predicted inten0.8
80sities much smaller than those observed. Alternative models propose the existence of low-lying
0.660
cool loops (15–17), but these static models cannot
be reconciled with the highly dynamic loops
0.440
we observe here. Guidance comes from threedimensional (3D) modeling: Low-lying, episod0.2
20ically heated loops that seldom or never reach00.0
coronal temperatures naturally arise in 3D
−100
−50
0
50
100
150 −150 −100
−50
0
50
100
150
models and predict (18–20) highly dynamic
km/s
km/s
spectraloriginatingin thewith
lower
transition
Fig. 4.lines
Episodic
structuresproperties
similar to UFS loops arise in advanced 3D numericalarc sectime [1000 s]

Intensity [DN]arc sec

Frequency

Frequency

0.25

insight into the ot
coronal heating mture, height distrib
UFS loops provide
3D models of the cother observations aof these loops are
relation between th
field in the photosp
the outer atmosphe

REFERENCES AND N

1. J. Mariska, The Solar T
Astrophysics Series (C2. U. Feldman, On the un
atmosphere in the 30,
Astrophys. J. 275, 367
3. U. Feldman, On the un
atmosphere. II. The tem
Astrophys. J. 320, 426
4. U. Feldman, K. G. Widi
quiet solar upper atmo
temperature regime. A
doi: 10.1086/307682
5. U. Feldman, I. E. Damm
fine structures of the swith the chromospheredoi: 10.1086/322471

Though it is established that the magnetic
field plays a major role in the energetics ofthe bright corona, determining the detailsof the physical mechanisms that heat thesolar corona remains one of the outstanding open issues in astrophysics. Several physicalprocesses are candidates for heating the corona,
including dissipation of magnetic stresses via reconnection and dissipation of magnetohydrodynamic waves (1–3). In many heating models, theenergy release is characterized by small spatial
and temporal scales. For instance, in the “nanoflare” model, random photospheric motions leadto braiding or shearing of magnetic field lines

and to reconnection, which yields impulse-driven
1
Harvard-Smithsonian Center for Astrophysics, 60 Garden
heating of the coronal plasma (4, 5). Several
Street, Cambridge, MA 02138, USA. 2Lockheed Martin Solar andstatistical studies of large numbers of solar flares
revealed rapidly variable emission,consistent
determinationof the plasma properties (e.g.,
Astrophysics
Laboratory,with
3251 Hanover
Street, Org. A021S,(6–8) have suggested that the energy-release mech3
Building
Palo Alto,reconCA 94304, USA.
Instituteof Theoreticalcoronal nanoflares causedby252,
slipping
plasmaflows
at different temperatures) andanism that produces flares is likely
similarAstrophysics,
Oslo, Post
Office Boxprevents
1029, Blindern,nection(12). within
However, thelack ofUniversity
spectralof intherefore
a detailed study of the physN-0315, Oslo, Norway. 4NASA Goddard Space Flight Center,a large range from micro- to X-classflares.
Ifnaformation in Hi-C data precludes an accurate
ical processes at work.
5
Greenbelt, MD 20771, USA. Dipartimento di Fisica e Chimica,noflares behave as a scaled-down version of larger
Universita’ di Palermo and Istituto Nazionale di Astrofisica
flares, then particles accelerated in the corona by(INAF)/Osservatorio Astronomico di Palermo, Piazza delreconnection processes could play an important
Parlamento 1, 90134 Palermo, Italy. 6Bay Area Environmental
Research Institute 596 1st Street West, Sonoma, CA 95476,
role in the heating of plasma even in the absence
USA. 7W. W. Hansen Experimental Physics Laboratory, Stanfordof large flares. Hard x-ray observations of microUniversity, Stanford, CA 94305, USA. 8Department of Physics,
27
flares (E ~ 10 erg) in active regions reveal the
Montana State University, Bozeman, Post Office Box 173840,
presence of nonthermal particles (8, 9), but
Bozeman, MT 59717, USA.
24
*Corresponding author. E-mail: ptesta@cfa.harvard.edu
nanoflare-size events (E ~ 10 erg) are not cur-

�Science IRIS 特集号

The Interface Region Imaging Spectrograph(IRIS) (16), launched in June 2013, is exceptionally well suited to investigate the response of thelower atmosphere to heating events, as it provides

Testa+2014

ナノフレアによるコロナループの加熱の証拠→ 電子ビームによる下層の加熱により上昇流が発生・数十秒から数分の短時間現象

E YE IN G TH E S U N
Fig. 2. Temporal evolution of TR emission for observed brightenings is dominated by time scales on the order of 10 to 30 s. Left: IRIS 1400 Å SJI of aregion where most of the footpoint chromospheric and TR brightenings occur. A few locations are marked for which the temporal evolution of thedeterminationof theemission
plasmaproperties(e.g.,IRIS SJI observations
The Interface
RegionImaging
chromospheric,TR, and coronalis shown
in Fig. 3. Right:of the same
subregion
shown in Spectrographthe left panel, at different times
[see
also
(19)
and
figs.
S1
and
S2].
Times
are
in
seconds
from
the
start
of
the
IRIS
observations,2013-11-09
12:04:16UT.
plasma flows at different temperatures) and(IRIS) (16), launched in June 2013, is exception-

revealed rapidly variable emission, consistent withcoronal nanoflares caused by slipping reconnection (12). However, the lack of spectral intherefore prevents a detailed study of the physally well suited to investigate the response of theformation
in
Hi-C
data
precludes
anaccurate
ical
processes
atwork.
lower atmosphere to heating events, as it provides
Fig. 1. Imaging observations of moss brightenings associated with coronal loop
heating.
Coronal and TR images of active region 11890, on 2013-11-09.
Fig.
3. Temporal
for by TR emission, whereas the 94 Å images is dominated by
The IRIS slit-jaw image (SJI) in the 1400 Å passband, and the AIA/SDO 193 Å imageevolution
are dominatedhot coronal emission (19). The brightenings that we focus on occur around x = 40 toobserved
65 and y = 60 to 90 [see also movies S1 and S2; in (19), we also discussbrightenings of line
brightenings that occur around x ~ 55 and y = 20 to 25].
and passband
intensities and Si IV
SCIENCE sciencemag.org
17 OCTOBER 2014 • VOL 346 ISSUE 6207
1255724-1
spectra shows a wide
range of upper- andlower-TR response.
(A to D) For thelocations marked (A)to (D) in Fig. 2, weshow (left panel) the
intensities versus time(with the minimum
value of the timeseries subtracted) of
FUV and NUV emission
lines (Si IV 1402.77 Å,
C II 1335.71 Å, Mg II k
2796.4 Å) observedwith IRIS, and in AIA
EUV narrow bands(171 Å, 193 Å); also
shown are the Si IV
spectra versus time(right panel), where
the wavelengths havebeen converted to
Doppler velocities (thedotted line marksv = 0). Negative

nsitionnarrow bright network lanes of this interface region. These jets have lifetimes of 20 toof convection
cells withthe pho-bright
80 seconds and widthsof ≤300 kilometers.
Theysimilar
originatesizes
from insmall-scaleregions, often precededby footpoint
brightenings
and accompanied
by holestransverse 466
tosphere.
Previous
observationsof coronalwaves with amplitudes of ~20 kilometers per second. Many jets reach temperatureswith the Solar Ultraviolet Measurements of Emittedof at least ~105 kelvin and constitute an important element of the transition region(SUMER)(4) ofonboardstructures. They areRadiation
likely an intermittent
butinstrument
persistent source
mass andthe
energy 464
for the solar wind. Solar and Heliospheric Observatory (SOHO) re-

Solar-Y (arc sec)

�Science IRIS 特集号

the solar wind (8). It has been difficult to test this
proposed idea without direct imaging of such TR
upflows on the solar disk, although moderateresolution observations have revealed signatures

Tian+2014

1

Harvard-Smithsonian Center for Astrophysics, 60 Garden
Street, Cambridge, MA 02138, USA. 2Lockheed Martin Solar andAstrophysics Laboratory, 3251 Hanover Street, Organization
A021S, Building 252, Palo Alto, CA 94304, USA. 3Max Planck
Institute for Solar System Research, Justus-von-Liebig-Weg 3,37077 Göttingen, Germany. 4Bay Area Environmental ResearchInstitute, 596 1st Street West, Sonoma, CA 95476, USA.
5
Department of Physics, Montana State University, Post OfficeBox 173840, Bozeman, MT 59717, USA. 6Institute of Theoretical
Astrophysics, University of Oslo, Post Office Box 1029, Blindern,0315 Oslo, Norway. 7High Altitude Observatory, National Center
for Atmospheric Research, Post Office Box 3000, Boulder, CO
762
766 768 770 772 774
80307, 764
USA.

462
vealedblue shiftsof 5 precise
to 10 origin
km s−1
forwithin coronal
2,4
he Suncontinuously
emitsDoppler
ionized particles
identifying
sites. Martínez-Sykora,
emissionlines
in theupperhigh-resolution
TR (5). They
into1interplanetary
space in theformformedof theholes
requiresobservations of
S. Saar,1 P. Testa,
solar wind. A challenging
investigationhas
thechromosphereandtransition
460 awere interpreted as signatures of the nascentregion (TR),2
now carried on for almost 50 years to uncomplex interface between the relatively cool phort,2 T. D. Tarbell,wind
by funnellike
hot corona (~106 K).
derstand wheresolar
the solar
windguided
originates
tosphere (~6 magnetic
× 103 K) and strucTheand energy
and how it is accelerated(1, originating
2). Dark regionsfrom
intures
the mass
networks(6). that ultimately feed458the
solar
wind
must
pass
this region.
coronal images indicate Recent
the coronal
holes
that
analyses revealed weak blue throughwing enThe dominant emission features in this interare the commonly accepted large-scale source
762 764 766 768 770 772 774
hancements
profilesfaceof emission
lines
formed*Corresponding author. E-mail: hui.tian@cfa.harvard.edu
region are the
network
structures that apsolar wind. in
However,romosphereregions
and of the high-speed
Solar-X (arc sec)
Solar-X (arc sec)
in
the
TR
(7,
8).
These
weak
enhancements
f the solar wind.480
Fig. 2. (A andSlit-jaw
B) Two 1330
unsharpmasked 1330 Å slit-jaw images showing the origin of network jets
Intensity
indicate the possible2014-01-23T07:28:39
presence of a plasma com- from small-scalebright regions in the network (movie S4). The red lines outline two jets.
al the prevalence ofponent flowing upward with speeds of 50 to 0.0
0.5
1.0
er second from the
−1
,
which
may
provideheated
mass
to
100
km
s
10
Slit-jaw 1330 Intensity
Peak Intensity
Line Width [km/s]
lifetimes of 20 to
the
solar
wind
(8).
It
has
been
difficult
to
test
this
8
all-scale bright 470
0.00
0.50
1.00 0.00
0.80
1.60 8.0
20.0
32.0
proposedidea
without
directimagingofsuchTRby transverse
6
2014-01-23T07:33:55
upflows on the solar disk, although moderate500
ch temperatures4
resolution observations have revealed signatures
ansition regionT

jaw imaging observations with the 1400, 1330, and
2796 Å filters [see the supplementary materials(SM)] unambiguously reveal the prevalence ofsmall-scale jetlike emission features from thebright networks (figs. S1 to S3 and movies S1 andS2). These three filters sample emission from the
Si IV, C II, and Mg II ions, which are formed attemperatures of ~105 K, ~3 × 104 K, and ~104 K,
respectively. These network jets usually show fast
upward motion with no obvious downward component. Although these jets are more easily seenin coronal holes located near the solar limb
(movies S1 to S5), they are clearly detected at
any location on the solar disk outside active regions (movie S6).
These network jets are best seen in 1330 Å
images. The jet widths are usually around ~300 km
and approach the instrument resolution limit,

1

Harvard-Smithsonian Center for Astrophysics, 60 Garden
0
Street, Cambridge, MA 02138, USA. 2Lockheed Martin Solar andAstrophysics Laboratory, 3251 Hanover Street, Organization
A021S, Building 252, Palo Alto, CA 94304, USA. 3Max Planck 12
Institute for Solar System Research, Justus-von-Liebig-Weg 3, 1037077 Göttingen, Germany. 4Bay Area Environmental Research
8
Institute, 596 1st Street West, Sonoma, CA 95476, USA.
5
Department of Physics, Montana State University, Post Office 6
Box 173840, Bozeman, MT 59717, USA. 6Institute of Theoretical 4
Astrophysics, University of Oslo, Post Office Box 1029, Blindern,0315 Oslo, Norway. 7High Altitude Observatory, National Center 2
for Atmospheric Research, Post Office Box 3000, Boulder, CO 00770 USA.
780
790
800
80307,
Solar-X(arcE-mail:sec) hui.tian@cfa.harvard.edu*Corresponding
author.
Number of jets

in sites within coronal450
olution observationsoftransition region (TR), a
en the relatively cool phond hot corona (~106 K).
440
hat ultimately feed therough this region.
on features in this interwork structures that ap-

20

speed = 206 km/s

100

1

200

300

400

a=155.1
s=62.20480

20a=49.8s=23.8

10
5

460

100

300

Speed (km/s)

0

50

100

150

100

Line profile at location 2

Lifetime (second)

440 17 OCTOBER 2014 • VOL 346 ISSUE 6207

2

12

2

10

12557 11-1

8CountsDistance (arc sec)

50

0

200

420

6

2

4
790 795 800 805 810 815

2
0

8

0

4

6

10

-50

Velocity (km/s)

8

12

20

0
-100

10

speed = 206 km/s
0of jets

1.0

30

10

2

0.5

40

15

Slit-jaw 1330 Intensity

SCIENCE sciencemag.org

1Time after 2014-01-23T07:25:06 (second)Fig. 1. Examples of network jets. (A) An unsharp masked (SM) 1330 Å slit-jaw image (movie S3). The dashed line marks the path of a jet. (B) Space-time plot
for the jet marked in (A). (C) Distributions of the apparent speeds and lifetimes for 63 jets. The average (a) and standard deviation (s) values are also shown.

0.0

1Line profile at location 1

Countsss and energy 460

Solar-Y (arc sec)

Solar-Y (arc sec)

Distance (arc sec)

ネットワークから高速のジェットが発生→ これまでに観測されたものよりずっと速い（250 km/s）・ネットワーク輝点の増光後に発生

100

200

300

400

Time after 2014-01-23T07:25:06 (second)
a=155.1
s=62.2020
15

a=49.8s=23.8

Solar-X (arc sec)

20

25

30

Time (min)

35

20

25

30

Time (min)

35

0
-100

-50

0

50

100

Velocity (km/s)

Fig. 3. Signatures of network jets in Si IV 1393.77 Å line profiles (movie S5). (A) Unsharp masked 1330 Å slit-jaw image taken at 07:33:55 UT on 23
January 2014. (B and C) Temporal evolution of the intensity and line width along the slit from a Gaussian fit of Si IV line profiles. The vertical line indicates theslit location in (A) and time of 07:33:55 UT in (B) and (C). (D and E) Observed line profiles (black) at the two locations indicated by the arrows. Red lines are theGaussian fits.

12557 11-2

17 OCTOBER 2014 • VOL 346 ISSUE 6207

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Filament%erupPon%Harra%et%al.%2014
The Astrophysical Journal, 792:93 (7pp), 2014 September 10

strophysical Journal, 792:93 (7pp), 2014 September 10

Harra et al.

Harra et al.

10:44:55UT

10:46:55 UT

10:51:55 UT−50

Y (arcsecs)

−100
−150
−200
−250
−300
−1100

−1000
−900
X (arcsecs)

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−1000
−900
X (arcsecs)

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Y-slit (arcsecs)

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-200
-220
-240
-260

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-200
-210

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X (arcsecs)-899

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Start Time (26-Oct-13 10:14:57)

Mg II k average

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-180
-200

−1000
−900
X (arcsecs)

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-300
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-180
Figure 5. In the left-hand column, the IRIS slit-jaw images in 2796 Å are shown at 10:45 UT, then at 10:50 UT (when the eruption begins to push the loops), at
10:52 UT, and at 10: 54 UT. In the images, the cool material that forms the core of the eruption is seen at the top right of the image. The hotter front is not seen at theseEIT%Movie:%Expanding%LoopsImada%et%al.,%PASJ%in%press

Fe XII 195.12 A
Intensity
VelocityLine Width
12 December 2006 UT 01:07

-50
-150

Solar-Y (arcsec)

-250

-200

0

-50

200

-200

0

200

-200

12 December 2006 UT 19:07

0

200

-150
-250
0

200

-50

400

0

200

400

0

13 December 2006 UT 01:12

200

400

-150
-250

0

30

200

400

Velocity [km/s]
15
-15
0

0

200
400
Solar-X (arcsec)
-30

40

0

200

Line Width [km/s]
60
100
80

400
120

EIS%observaPon%before%the%large%ﬂare%(2006%Dec%13).%Flare%occurred%02:13%UT.%The%leZ/middle/
Fig. 7. The EIS .
right%panel%shows%the%FeXII%intensity/velocity/line%width,%respecPvely.%BlueshiZ%component%
was%observed%5%hours%before%the%ﬂare%onset.

Imada%et%al.%2014%PASJ
1st Stage:
Outer AR (Low Density) Expansion

2nd Stage:
Inner AR Expansion

3rd Stage:
Main FlareFirst%stage%%(a%day%before)%:%%
low%dense%outer%shell%expands%
Origin%of%non$thermal%vel.%%Spectroscopic%obs%only%%
Second%stage%%(a%few%hours%before):%%high%dense%core%region%expands%
Imaging%obs%also%
Flare%occur!%
13%Dec%2006%event%is%opposite%%polarity%event.%

Fig. 11. Schematic illustration of the observational results. The arrows represent the expansion of the coronal loops.

Hinode/EIS%observaPon
2014l081619ƜǬĞ%QƆƀ÷ěƗĒèø¼ƌƑ%ǫǪ(ƿǧǊǥǄ

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ƮÆĂƅƙƜăÅƌƖƮƜƒưƁƄǮ%シンポジウム「スペース太陽物理学の将来展望」

21-22日10月2014

Small-scale structure & dynamics in thecorona—Insights from coronal rain observations
with Hinode & IRIS
Patrick Antolin
国立天文台

Normalized Intensity

thermal width of σT = 0.32 MK.
Nevertheless, the thermal width of even t
than the majority of the loops in the sample
1.2
top panel of Figure 4, for example, has a therm
1.0
HiC
0.04 MK. From our EM analysis of the com
0.8
AIA
all the loops detectable with AIA, we foun
0.6
σT !0.32 MK. These results are comparable
measurements by EIS (Warren et al. 2008),
0.4
relatively more loops with a broad tempera
0.2
reflecting
the larger uncertainties in the AIA• High Reynolds numbers0.0in corona -> need for fine structure
for effective
results do, however, support the idea that01000 2000 3000 4000 5000
narrow temperature distributions and are com
Distance(km)
dissipation of waves, magnetic fields(currents)
and turbulence
few magnetic threads.
Recently, Peter et al. (2013) analyzed se
How is the magnetic field organised in the solar corona? observedby Hi-C and found that they have s
intensity profiles. They suggested that theymonolithic structures
or arecomposed of m
• Coronal loops evolve on timescales much longer than a radiative
coolingtime
strands. With our larger sampling of the data
(Reale & Peres 2000; Warren et al. 2003, Brooks et al. 2012)
some cases of long loops that show evidenc
when imaged by Hi-C. The loop shown inof Figure 4 2006,
is the best
example. The result➙ ensemble of unresolved independently heated strands? (KlimchukReale
are not atypical. The widths of loops like th
1.2
2010, Brooks et al. 2013, Peter et al. 2013, Cirtain et al. 2013, Winebarger
et al.
2013)3. We do agree with t
distributionof Figure
1.0
HiC
Peter et al. (2013) that many of the narrow fe
0.8
AIA
field of view appear to be relatively shor
nanoflares?
(Parker
1988, VeksteinHi-C
2009)• Strands ⟷ current sheets,
The instrument performance achieved by
0.6
demonstrates that 100 km spatial scales w
0.4• Clear tendency for finer widths
at higher spatial resolutionobservable with future coronal spectromet
0.2
planned for Solar-C (Teriaca et al. 2012). The
0.0
here are encouraging to the view that such in
01000 2000
3000 4000 5000
widths
TRACE ~ 1000 Mm
able to measure the true plasma properties
Distance (km)~ 100 kmand provide realistic constraints for coronal
If confirmed by these instruments, then our
measurements already represent the true spatia
structures. Theoretical models need to explain
structured on these scales, and why the tempera
are narrow over scales of hundreds of km.
Normalized Intensity

Importance of small-scale structure?

SST/ CRISP H!

ed Intensity

We acknowledge the High resolution Co
SDO
Hi-C
strument team for making the flight data p
MSFC/NASA led the mission and partners inc
1.2
nian Astrophysical Observatory in Cambridg(Aschwanden & Nightingale 2005,
1.0 Reale 2010, Peter et
(Lin
et
al.
2005,Martin’s
Antolin& Rouppe
HiC
Lockheed
Solar AstrophysicalLabora0.8
AIA
California; the University of Central Lancash
al. 2010, Brooks et al. 2012,
2013)
van
der Voort 2012)
0.6
England; and the Lebedev Physical Institut

-200

-210

10:50 @ y=-200"

60

10
0

Intensity

-205

20
km/s

Y-slit (arcsecs)

Wavelength (A)
-195

-10
-20

40

Tracer of coronal turbulence
-215

10:25

10:30
10:35
10:40
10:45
Start Time (26-Oct-13 10:21:26)

10:50

-210

0.4

compression?

•
•
•
•
•

Figure 5. In the left-hand column, the IRIS slit-jaw images in 2796 Å are shown at 10:45 UT, then at 10:50 UT (when the eruption begins to push the loops), at
10:52 UT, and at 10: 54 UT. In the images, the cool material that forms the core of the eruption is seen at the top right of the image. The hotter front is not seen at thesewavelengths. The right-hand column shows sample spectra of Mg ii along the slit at each time. The spectra shown on the top right are at the same times as the images.
The plasma in the loops show multiple features. At 10:50, once the eruption pushes the loops, the spectra are less complex with narrower line profiles, indicating
simpler dynamics. At 10:52, the plasma shows a strong blueshift, and at 10:54, the plasma then becomes redshifted (please note that the y axis has changed in order totrack this feature). These strong flows are seen while the loops are being pushed downward. The animation shows the spectra changing with time following the eruption.
(An animation and a color version of this figure are available in the online journal.)

2804

Figure 7. Sample Mg ii spectra at y = −200"" before the filament eruption (top)and after the filament eruption (bottom).

0.3
0.2
0.1
0.0

10:25
Harra et al.

2798 2800 2802
Wavelength (A)Mg II k/h ratio
10:30
10:35
10:40
10:45
Start Time (26-Oct-13 10:21:26)

10:50

Figure 6. Two-component fit applied for the Mg ii data. These figures show theMg ii slit data with time for the main fitted component—intensity is shown in
the top panel, Doppler velocity in the middle panel, and line width in the bottom
panel.
(A color version of this figure is available in the online journal.)-195

observed following eruptions (Nakariakov et al. (1999)). Thisis the first time this has been observed.
We determine the Mg ii h/k ratio along the slit with time todetermine if this changes during the impact of the filament.
Examples of the Mg ii h and k spectral lines are shown in
Figure 7 spatially located in the center of the loop structure.
At 10:35 UT before the filament eruption, the profiles shownno central reversal. In addition, the spectra are non-Gaussianbroadened profiles compared with the spectrum of the isolated
blob (Figure 2). There is evidence of a red-wing component.
The ratio of the h/k line at this time is 1.5—similar to that of
the isolated blob. At 10:50 UT just as the filament impacts theseloops, the profiles again show complexity with stronger bluewing components. The ratio of this stage has increased to 1.9.
To show the temporal and spatial evolution of the intensity ratio,
we determine this value in all the pixels that have statisticallysignificant intensity values. Figure 8 shows the stack plot of
the ratio. Most of the pixels have a ratio around 1.6 before thefilament impacts. The ratio increases to two and above following
the impact. At no stage does the ratio reach a value of four, which
would indicate that radiation is the dominant process.
We are seeing significant changes in these high-lying loops
following the impact of the filament—both in terms of thecomplex dynamics within the loops and in terms of the line
intensity ratios.

-200

-205

-210

Mg k/Mg h-205

Y-slit (arcsecs)

before impact oferuption
during impact

The Astrophysical Journal, 792:93 (7pp), 2014 September 10

2796

-200

-215

coronal rain

0

-195

Angstrom

Y-slit (arcsecs)

Line width

20

2.5
2.0
1.5
1.0
0.5
0.0

-215
10:25 10:30 10:35 10:40 10:45 10:50
Start Time (26-Oct-13 10:21:26)
Figure 8. Stack plot of the Mg k/h ratio.
(A color version of this figure is available in the online journal.)

Complex spectral evolution, mostly due todynamics
During impact of eruption: complexity
momentarily reduced (compression?).
h/k ratio increases 1.5->2.
Significant plasma reorganisation occurs
within loops during eruption
Rain can serve as a probe for turbulence
3. DISCUSSION

These observations show coronal rain above the limb observed for the first time in Mg ii by the IRIS spacecraft.
The line profiles show that the central reversal is gone. The
complexity seen in the spectral lines can then be assumed to be

6

Harra+ 2014

Supersonic heating events
IRIS SJI + SGt) and SDO images at the same time (15:16 UT on 2013 August 30). Faint loop parts can be seen in the
AIA
304193 (white arrows). The bright dots visible in IRIS images are located atbright loops in AIA
171 andarrows). See movie.

eeds (below 50 km/s) thanof projected to line-of-sightis however similar to that
ations of rain falling almost
In that case, the projected
mponent. Using Equations2012) adapted to an on-disk
that the observed projected
onds to an angle of ∼5–20
respect to the loop baseline. Our observations also re(2013), although on a much

• High red shifts (up to 200 km/s) observed in (E)UV brightpoints50-100 km/s upflows
4.2. accompanied
Similarity to Flareby
Ribbonsp
Thermal instability appears to be the same for postcs =(Foukal
RT⇡ 50 km/s• loops
Strongly
supersonic:
flare
and non-flaring
active regions
1978;

Schmieder et al. 1995). The observation on 2013 July
Brightdots
arefootpoints
ofAugust
thermally
unstable coronal
19 •showedslightly
more
activitythan 201330,
but there
were only
microflares,inthe304strongestof whichto Doppler shifts
loops:
downflows
correlated
occurred toward the end of the observations while theemission dots were fading (see online movie). Coronalappear(within umbra or
rain•wasMay
visiblein AIA anywhereimages duringin
thesunspot
phase of thesimple ribbon(before 7.40
UT).forThen,a jet-likepenumbra)& lastmorethan activ2 hours
ity ejected plasma, which can be seen falling back down
Kleint+ 2014
in AIA 304, additionally to the coronal rain, and whichcoincides with the time when the ribbon changes shape.

Hinode-IRIS co-observing campaign

Hinode/SOT
FOV

Antolin+ (2014a, in prep.)

Multi-temperature structure
Hinode/SOT + IRIS/SJI + SDO/AIA
Green: SJI 1400,
log T ~ 4.8 - 5

Blue: SJI 2796,
log T ~ 4 - 4.2

Red: SOT Ca II
H, log T ~ 4 - 4.2• Chromospheric cool, compact and dense cores surrounded by diffuse warmer plasma?
• Emission time delays from TR to chromospheric emission observed. Very fast evolution:catastrophic cooling• Exception: SDO 171 & 304: low T contribution? (In contrast with Kamio+2011,Peter+ 2011)

Multi-stranded structure: tip of the iceberg?

• Structural differences withtemperature• Co-spatial chromospheric & TR
The Astrophysical Journal Letters, 771:L29 (6pp), 2013 July 10
emission
Numerical model• Peak of distribution at very low(a)spatial scales: tip of the iceberg
Fang et al.
scenario? (Fang+’13, Scullion+’14)(2013)• Less uniformity for lengths• Effect from temperature?
Interaction with magnetic field?
(Sharma+2010)
Observational Resolution ΣN=1103
Numerical Resolution ΣN=4146

140

120

120

100

100

80

80

N

N

140

60

60

40

40

20

20

0

0

0.5

1

1.5

2

0

0

0.5

Width (Mm)

Figure 4. Panels (a) and (b) show the distribution functions of width and length, respEstimates of non-thermal line broadening

Semi-automatic detection of rain/prominence plasma for statistical analysis
Antolin+ (2014b, in prep.)Estimates of non-thermal line broadening
FWHM ="✓p
2 2 log 2

0

c

◆2 ✓

2kB T
2
+ ⇠nthm• Mostly single emission peaks• Gaussian-like distribution of
non-thermal broadening withvalues < 25 km/s and a peak
≲10 km/s (consistent with log(T ) = 4.8
Si IV
Parenti & Vial 2007)
log(TC II ) = 4.3• Increase of non-thermal
broadening with height(consistent with Okamoto+2014, submitted)log(TMg II ) = 4.

Height dependence◆2+ ⇠inst# 12

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�

Mg II の例
〜 リム外のプロミネンス 〜岡本�Mg II の例1
Mg II kSOT 画像上部は視線方向の重なりが少ない(20秒間隔)時間

波長(k/h 比)Mg II h�Mg II の例2
Mg II kSOT 画像重なりは比較的少ないものの、
central reversal が見られる
これは multi-component
か、optical effect か？

(20秒間隔)時間

波長(k/h 比)Mg II h�Mg II の例3
Mg II kSOT 画像
複数のスレッドが視線方向に重なっていると考えられるが、プロファイルはシンプル(20秒間隔)

ライン幅も狭い時間

シングルのものが小 k/h かというと、そうでもない

波長(k/h 比)Mg II h�Mg II の例4
Mg II kSOT 画像重なりは多いcentral reversal が見られる(20秒間隔)

ライン幅は広い時間

どの効果が一番効いているのか？

波長(k/h 比)Mg II h�スレッドの振動

Okamoto+2014, submitted

背景： カルシウム線 （ひので） 画像の x-t プロット
黄点： 微細構造 （ひので） の中心位置
赤点： マグネシウム線 （IRIS） の視線速度

マグネシウム線は不透明度が高い

b S2

a S1

500

5

0

0

and Ca II H (line center) for a cross-section of the tube along the tube’s centre. -500
The yellow curves correspond to contour lines for different values of the-5

Height along slit (km)10Supplementary Figure S2 : Source function in Mg II k 279.635 nm (left panel)

微細構造表面のoptical thickness (𝜏, stated in the legends) as seen from the top (as indicated -1,000
運動を反映
by the arrow), corresponding to a LOS angle of 0°.

-10
0

100

200

300

400

tau
0.05

1.00

0.10

3.00

0.25

6.00

0.40

0

−2

Height along slit (km)y/R

2

Ca II H

0.30

0
x/R

2

−2

0
x/R

400

300

500

鉛直振動と視線方向の振動が 180度ずれている10

1,000
500

5

0

0

-500

-5

-1,000

-10

0
−2

200

100

共鳴吸収

300
200
Time (sec)

400

0

100

200

300

400

Time (sec)2resonant absorption(e.g., Sakurai+1991, Goossens+2002)

Line-of-sight velocity (km/s)

tau

100

d S4

c S3

Mg II k0

Line-of-sight velocity (km/s)

1,000�スレッドの振動

Okamoto+2014, submitted

共鳴吸収

b S2

Height along slit (km)

a S1
1,000
500
0
-500

-1,000
0

100

200

300

400

Height along slit (km)

c S3

0

d S4

1,000
500

鉛直振動と視線方向の振動が180度ずれている
0-500-1,000●�微細磁束管の密度成層により、表面で振動周期が減少
0
400
300
100
200
● 微細構造表面で Kelvin-Helmholtz
不安定性を誘発Time (sec)

振動のエネルギー�

表面運動のエネルギー

微小な電流層を生成し、振動のエネルギーを散逸
Antolin+2014

10シンポジウム「スペース太陽物理学の将来展望」

21-22日10月2014

Alfvénic wave detection in the solar
atmosphere—
Possibilities for Solar-C
Patrick Antolin
国立天文台

Observations of Alfvénic MHD waves
in the solar atmosphere
Transverse MHD waves ubiquitous in the solar atmosphere(Tomczyk+ 2007, Okamoto+ 2007, Terradas+ 2008, De Pontieu+ 2007, Lin 2011, McIntosh+ 2011, Morton+ 2011,
Antolin & Verwichte 2011, Okamoto & De Pontieu 2012, Hillier+ 2013, Schmieder+ 2013, Morton & McLaughlin
2014, De Pontieu+ 2014)➙ small amplitude (~km/s), periods of few min

Coronal loops &prominencesthat the heating we observed is substantial. Someof these events are also visible in coronal passbands (movie S7), although it is not yet clear how
much plasma is heated to coronal temperaturesin association with these heating events (12).
Our observations of twist that permeates thechromosphere and the TR of the Sun expand ona picture that has recently emerged that draws
together disparate observations of twist in the

Chromosphere

Fig. 2. Spatiotemporal
properties of twist.
C
D
A
B
Rapidly evolving
twisting motions areapparent as shortlived, bright features inthe blue and red wings(e.g., around T50 km/s)of the chromospheric
Mg II h 2803-Å spectralline (C and D) in
regions of inclined field[positions 2 and 3
shown in (A), Si IV
E
F
G
1403-Å SJI]. Thesemotions are in contrast
to the acoustic shock–dominated spectral
profiles (B) in position
1 in plage (where field
and line of sight are
more aligned,
preventing visibility oftwist) that evolve ontime scales of several
minutes (25). The
spatial pattern of thebright Si IV feature around –190″, –207″ [position 2 in (A)] is associated with short-lived twisting motions that are visible in Si IV [80,000 K (G)], in Mg II h
[10,000 K (F)], and faintly in Ha [<10,000 K (E); see movies S5 to S7]. Velocities of order 50 km/s are reached in this twisting feature. Typical velocities
are lower (10 to 30 km/s) with visibility in these various passbands variable; Mg II h showed excellent visibility in most cases. See supplementarymaterials for other examples (figs. S7 and S8) and movies S5 to S19 and S21.

SST + IRIS

Tomczyk et al. 2007SDO

SCIENCE sciencemag.org

Antolin & Verwichte 2011

t=0 s

t=15 st=30 st=44 st=59 st=74 st=89 st=104 s6

Y [Mm]

CoMP

Fig. 3. Temporal
evolution of twist andassociated heating.
SST Ha dopplergrams
at T46 km/s (bottom;
black is blue-shifted;
white is red-shifted)show how quickly
chromospheric twist
propagates along
elongated features ontime scales of less than
1 min. Several of thesetwisted features areassociated with TR
signals (top) as
observed with the IRIS
SJIs that are dominatedby Si IV lines (~80,000 K).
Movies S8 to S10provide further examplesof the dynamic nature of
the twist and associatedheating. t indicates time.

4
2
0
6

Y [Mm]

Okamoto et al. 2007

McIntosh et al. 2011

dynamic linear features (C II 1335 Å and Si IV
1402 Å); this emission indicates heating to at least
TR temperatures (20,000 to 80,000 K). In quiet
Sun regions, these SJI features appear to be theTR counterparts of so-called rapid-blue-shifted
and rapid-red-shifted events in the chromosphere,
the disk counterparts of spicules (26). Many of the
TR features IRIS observes in quiet Sun regions
are associated with twisted features, indicating

Hinode/SOT

Hinode/SOT

Large-scale corona

The torsional time scales stand in sharp contrast
to the longer time scales (several minutes) andsmaller line-of-sight velocities in plage regions(Fig. 2B), where the line of sight is more alignedwith the magnetic field and the line profiles are
dominated by magnetoacoustic shocks (25).
SJIs taken with IRIS (Fig. 3 and movies S8 toS10) also show how the propagating torsionalmotions are often associated with bright, highly

4
2
0
0

1

2

3

0

1

2

3

0

1

2

3

0

1

2

3

0

1

2

3

0

1

2

3

0

1

2

3

0

1

2

3

X [Mm]De Pontieu et al. 201417 OCTOBER 2014 • VOL 346 ISSUE 6207

1255732-3

Significance of MHD wave identification
Alfvénic waves can deliver sufficient energy fluxes in order to heatand maintain a corona
Uchida & Kaburaki 1974, Wenzel 1974, Hollweg+ 1982, Poedts+ 1989, Ruderman+ 1997, Kudoh & Shibata 1999,
Suzuki & Inutsuka 2006, Cranmer 2007, Antolin & Shibata 2010, Matsumoto & Suzuki 2014

• Energetically significant: Alfvénic➡ Trapped / non trapped; Torsional, Kink, Sausage / Fast, Alfvén• Energetics & dynamics: each mode carries different amounts ofenergy & involves different plasma motions• MHD Seismology: each mode has different dependence on thephysical properties of medium

Requirements for wave identification?

Transverse MHD wave
(kink mode)No. 2, 2008

Antolin+2014

NONLINEAR INSTABILITY

1. Resonant absorption
Coronal Waves and Oscillations

Nakariakov+1999‣‣

Conversion of a global kink oscillation into local Alfvénic
oscillations (Hollweg & Yang 1988, Goossens+1992, 2002, Ruderman &Roberts 2002, Terradas+2010, Pascoe+2010).

Very successful in explaining observed fast damping(Nakariakov+1999, Verth+ 2010, Arregui & Ballester 2011)

Figure 10: The temporal evolution of the loop displacement as an average coordinat
position for four neighbouring, perpendicular cuts through the loop apex (diamondsbars (±0.5 pixels), starting at 13:13:51 UT on 14 July 1998. The solid curve isthe function A sin(!t + ) exp( t) with A = 2030 ± 580 km, ! = 1.47 ± 0.05 rad= 0.069 ± 0.013 min 1 , corresponding to a period and e-folding decay time of P = 4
and ⌧ = 14.5 ± 2.7 min, respectively (from Nakariakov et al., 1999).

Terradas+ 2008

main di erence of this polarisation from the horizontal one is that in the curved loop tpolarised kink mode changes the length of the loop. Consequently, as the mass in the
be conserved, this mode can have a significant compressible component. The fractio
perturbations associated with this mode were estimated as
2∆ ⇢
2∆ L
∆I
/
/
,
I
⇢
L

2. Kelvin-Helmholtz instability (transverse to main field)‣

where ∆ ⇢ is the density perturbation and ∆ L is the perturbation of the loop length
assumed that the oscillation does not deform the loop cross-section, i.e., ∆ a/a ⌧alternative is associated with bulk field-aligned flows through the loop footpoints, w
have observational confirmation.
Another possible polarisation of kink oscillations is when the loop moves in the sam
its apex displaces horizontally. This mode has not been identified yet.

Deformation of cross-section: strand-like structure in intensity
images (Antolin, Yokoyama, Van Doorsselaere 2014)
3.2

Fig. 1.—The top half of each panel displays the density (symmetric respect to y

L /2 for observations
different times. In this simulation the following parameters have been used:
Non-TRACE

6R] # [0, 10R]. Lengths are normalized to the loop radius R (typically of the order oModel: Uchimoto+. 1991; Karpen+ 1993, Ofman+ 1994; Ziegler & Ulmschneider 1997;
Terradas+ 2008, Soler+ 2010; Obs: Foullon+ 2011; Ofman & Thompson 2011, Berger+ 2010of 1000
km –s often), and erroneously
time to the Alfvéncrossing
time,waves
tA p R/
vAi. [This figurebeca
is aKink modes of coronalloops
calledAlfvén(erroneously,
Alfvén waves guided by coronal loops are the torsional modes, see Section 8) – can also
through modulation of the broadband gyrosynchrotron emission by the periodic var
largeandlongitudinallength scales
preventofthe
local magnetic fieldHowever,in flaringthe
loopsthrough periodic
variations
theinterpreDoppler shtationofequation
(1)
as
a
local
approximation
in
z
since
the
velocity
The optically thin gyrosynchrotron emission intensity If at a frequency
f can bshear
less than
2 v0 throughout
most formulaof the loop. Growth rates
with the use of (Dulkandis Marsh,1982)’s
approximated!1

are thus expected to be somewhat smaller, and instability will set✓◆

Spectral and imaging signatures of
the transverse MHD wave
Prominence model
Coronal model
Possible with
SOT+IRIS
Antolin+(2014)

0.1 R

0.3 R

0.5 R

1.0 R

Okamoto+(2014,sub.)Possible for
Solar-CTransverse motion observable with resolution of 0.2 R
Spectroscopic resolution of 0.5 R resolves azimuthal flow
➡ Strand-like structure in intensity images (scales ~ 0.1 R)➡ Phase difference between displacement in POS and LOS velocity due
to higher propagation speeds of Alfvénic waves in resonant layer

Importance of forward modelling:
opacity effects

Supplementary Figure S2 : Source function in Mg II k 2796.35 A (left panel)
and Ca II H (line center) for a cross-section of the tube along the tube’s centre.
The yellow curves correspond to contour lines for different values of theGoossens et al.

optical thickness (𝜏, stated in the legends) as seen from the top (as indicatedby the arrow).

3D MHD simulations + RH

Assuming thick cylinder(simplistic approach)(Okamoto+ 2014, submitted)

Transverse MHD wave

1

Torsional Alfvén wave

2

3

Ik
⇡ 1.5
Ih

2

1

KHI flow2

Dipole-like Alfvénic
flow in resonant layer

3(Goossens+2014)

Figure 9. Integrated intensity of the transverse kink mode (top window) andthe torsional mode (bottom window). In both windows, the intensity is shown
as a function of the Doppler shift (horizontal axis) and the position across thecylinder (vertical axis). The color scale of light pink to dark pink indicatesincreasing emission, and white is no emission. The top window contains several
panels corresponding to different viewing angles between the LOS and directionof classic kink transverse motion (indicated in the top left of each panel). The
contour lines in the top window are shown for 1%, 5%, and 20% of the maximum
intensity when the LOS angle is 90◦ relative to the direction of the classic kinkmotion.
(A color version of this figure is available in the online journal.)• Emission comes mostly from ring, mostly optically thin. The prominencecore is optically thick in Mg II h&k, not in Ca II H&K: advantage of Mg II
Properdynamics
only
reason •for thechange of Doppler
sign from the center
to the recovered with proper forward modellingboundary of the tube in Figure 6(d) was because CROT waschosen to be larger in magnitude and opposite in sign to CTR at輻射流体計算を用いた
太陽彩層研究について
飯島陽久
東大・地惑

2014年10月20-21日シンポジウム「スペース太陽物理学の将来展望」＠JAXA/ISAS 相模原キャンパス

内容•

数値計算におけるH, Ca, Mgのモデリング•

輻射流体計算と観測との比較•

現在の輻射流体計算における制限

彩層のダイナミクスに関わる元素

Vernazza et al. (1981)

-

水素の電離平衡までの時間スケー
ルは彩層で数分から数時間と長い。非平衡電離の効果はダイナミ

クスにもポストプロセスでの輻射

水素原子由来の輻射冷却率

輸送計算にも重要。

電離平衡までの時間(log10[sec])

水素の取り扱いについて-

Carlsson & Stein (2002)

HαやLyman線、コンティニュ
ウム等は彩層における輻射冷却に大きく寄与する。

Vernazza et al. (1981)

Ca, Mgの取り扱いについて-

CaやMgの取り扱いは状態

彩層における全輻射冷却率への寄与

方程式としてはそれほど大きな影響は無い。
-

しかし、Ca II K&H、Mg II
h&kは彩層の輻射冷却率において重要。

-

ポストプロセスの輻射輸送計算においては(水素の非平衡電離が考慮されていれば)統計平衡がかなり良い近似になる(Wedemeyer-Bohm
Vernazza et al. (1981)& Carlsson, 2012;
Leenaarts et al., 2013a)。

彩層輻射冷却のモデル化(Carlsson & Leenaarts, 2012)
ラインによる冷却率-

H, Mg, Caからの輻射冷却を温度、密度、光学的深さの関数としてモデル化。

-

静穏領域の1次元RHD計算から
フィッティング。

コロナ輻射による加熱率

実線: 3D-

コロナから輻射による彩層の加熱は1次元(カラムbyカラム)輻射輸送

破線: 1Dを解いてモデル化。
-

多次元性を考慮していないこと（局
所的には正しくないこと）、輻射
「加熱」の近似精度等に問題点。

Carlsson, 2010, NAOJスライド長さの最大値

スピキュール(Martinez-Sykora et al., 2009/2013)

速度の最大値-

左図: 横から見たHα強度

右図: スピキュールの長さと速度の相関。

スピキュールが短い。

左図のような背の高い構造は非常にまれ。また観測される
II型スピキュールより密度が小さい。

Ca II 854.2 nmスペクトル(Leenaarts et al., 2009)
観測: Wing

観測: Core
赤: 観測
黒: 計算

計算: Wing

計算: Core-

空間平均したライン幅は観測より狭い。グリッド数不足により解像されていない、細かいダイ

ナミックな構造の存在を示唆。
-

ラインウイングのイメージは観測とよく似ているが、ラインコ
図: CRISPによる観測とBifrost
による数値計算の比較。

アは細かい構造が消えている。

Hαフィブリル様構造(Leenaarts et al., 2012)

計算(Bifrost)-

3次元RMHD計算のHαでフィブリル様構造を初めて再現。

観測(SST)-

送計算では見えない)。
-

図: Hα線コアの強度。SST(左)とBifrostによる計算(右上)。

3次元散乱が重要(1次元輻射輸
フィブリルは観測よりまばらで長さも短い。また、ライン幅もずっと小さい。

Mg II h&kの形成(Leenaarts et al., 2012/2013ab)-

完全な3次元PRD計算はポストプロセスであっても難しい。

-

k1, k2についてはカラムbyカラ
ムの1次元PRD計算、k3に関し

1D: k2vては散乱の効果が強いので3次

1D: k3

元CRD計算が良い近似。
-

モデル大気において水素の非平衡電離が考慮されていれば、得られたTとNeを用いて統計平衡を仮定してMg IIのラインを計算
3D: k2v

3D: k3

図: CRDを仮定した場合の1Dと3Dの比較

してもかなり正確。

計算コストの問題•

Bifrostコードの1コア辺りの性能は3×10^4 grid*step/sec程度。
•

例えば500^3グリッド(24x24x15Mm^3, 48x48x30km^3)、時間幅が0.01秒の計算を1太陽時間行うには3000コアで2ヶ月かかる。相当大変な計算。

•

グリッドサイズへの制限•

粒状斑の分解(〜100km)、彩層のスケールハイトの分解(〜30km)、ローカルダイナモ(〜
10km)、遷移層の厚みの分解(〜数km)など。Ca II 8542等のライン幅不足(Leenaarts+,
2009)もグリッドサイズ不足が考えられる。

•

領域サイズへの制限••

より大きい対流構造の影響を取り入れようとすると下部境界を深くする必要がある。

磁場強度への制限•

黒点や活動領域付近の大気ではアルフベン速度が非常に早くなるため、シミュレーションの時間幅を強く制限される(10^-3秒以下)。また、計算が壊れやすくなる。

輻射流体計算の課題••

観測との比較
•静穏領域 Ca II 854.2 nm: ライン幅が狭い。グリッド数不足が推測される。

•

静穏領域 Hα: フィブリル様構造が再現され始めている。密度や長さなどは足りない。

•

活動領域、Ca II H&K, Mg II h&k等: （多次元の）詳細な比較研究なし。

計算コスト
•

•

すでに計算コストはかなり大きい。これ以上大きい計算や活動領域の計算は難しい。

考慮されていない物理過程•

Hall効果、Ambipolar拡散•

Ca, Mg, Heの非平衡電離•

3次元効果を取り入れた彩層の輻射冷却率

Appendix

数値計算による彩層スペクトル研究•

Leenaarts et al., 2009: Ca II IRT (SSTとの比較あり)•

注：モデル大気はLTE(水素の非平衡電離は考慮されていない)。

•

Leenaarts et al., 2012: Hα (比較あり)•

Leenaarts et al., 2013ab: Mg II h&k•

Cruz-Rodriguez et al., 2013: Ca II IRT (SSTとの比較あり)••

彩層加熱とCa II 8542

Riethmuller et al., 2013: Ca II H, Mg II k (Sunriseの観測。Mg II kは計算との比較あり)•

Leenaarts+13と比較。ただし、SuFIフィルターの幅が大きいため詳細な比較には至っ
ていない。

•

Leenaarts et al., 2014: Ca II IRT•

同位体の効果未考慮の物理過程: ヘリウムの非平衡電離-

ヘリウムの電離平衡までの時間スケールは数100秒。

Leenaarts et al. (2011)-

1万度付近でHe IIがHe Iに変わる時に潜熱を開放するため、温度が変わりにくくなる。

-

Golding et al. (2014)

この効果は電離非平衡を取り入れると小さくなる。

未考慮の物理過程: Hall効果, Ambipolar拡散
Martinez-Sykora et al. (2012)•

磁場が強く温度が低い場所で、磁場のダイナミクスやエネルギー的に重要になる。

未考慮の物理過程: 多次元性を取り入れた輻射冷却率--

Escape probabilityの形でモデル化されたMg II線の輻射冷却は、平均成層は良く再現するが、局所的にはかなり差異がある。
また、輻射による加熱についてはモデリングが難しく、考慮されていない。

-

1次元でモデル化されたコロナからの輻射の吸収も、3次元の場合はよりスムーズな加熱になっている。