光球-彩層間の波動現象と ALMAによる多層間結合撮像

光球-彩層間の波動現象と
ALMAによる多層間結合撮像
加藤 成晃（NAOJ）
研究協力者
Mats Carlsson (ITA, Univ. of Oslo)
Jorrit Leenaarts (ITA, Univ. of Oslo)
概要
1. リアリスティック太陽大気モデリングとは
2. 光球-彩層をつなぐ磁束管での波動現象
3. ALMAによる多層間結合撮像への期待
ALMA太陽観測WS 2012/09/04
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1
磁束管で自発的に生じる
磁気流体波とは？
Corona
Chromosphere
Dissipation
Propagation
Photosphere
Convection
zone
β=1
τ=1
Generation
Hinode/G-band
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2
磁束管で自発的に生じる
磁気流体波とは？
Corona
Chromosphere
Longitudinal
wave
Dissipation
Propagation
Photosphere
Convection
zone
Torsional
wave
β=1
τ=1
Generation
Pumping Buffeting Swirling
Hinode/G-band
ALMA太陽観測WS 2012/09/04
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Transverse
wave
2
Why we need realistic modelings?
18
16
14
8
7
log10 Te [K]
6
5
Corona
10
Chromosphere
12
Photosphere
log10 ne [1/cc]
Hinode/Ca II H
4
THE HARVARD-SMITHSONIAN REFERENCE ATMOSPHERE
3
ALMA太陽観測WS 2012/09/04
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349
Temperature
this new reference model. The
details of the
1 model are given in several
10tables, which
Z [Mm]
appear at the end of the paper.minimum
We wish to emphasize that the HSRA, like all models that have preceded it, should
not be interpreted as an accurate description of the true solar atmosphere. The model
describes an idealized plane-parallel homogeneous atmosphere in hydrostatic equilibrium; however, it is well known that in the Sun, and especially in the chromosphere,
perturbations due to magnetic and hydrodynamic effects grossly distort the local
structure from its mean configuration.
In addition, we expect this model, like the BCA and the Utrecht Reference Photosphere that preceded it, to be relatively short-lived. For example, new observations
and analyses3in the ultraviolet and infrared can be counted on to improve the mean
structure in the chromospheric layers, and to shed light on inhomogeneities. Never-
Toward Realistic
Modelings of Solar Atmosphere
Heat conduction along
field lines from corona
Basic equations
Radiative Transfer
- LTE + Non-LTE
- Opacity
✓
✓
Continuum
Lines
- Ionization
✓
✓
Collisional excitation
Photo-ionization
Magneto-convection
- Radiative cooling
- EOS
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Solar “ingredients”:
4
Waves inside Flux Tube
Overview
Analytical Solutions of Wave Equations
- Defouw 1976; Roberts & Webb 1978; Spruit & Zweibel 1979: Longitudinal waves
- Roberts 1981; Rae & Roberts 1982; Spruit 1982,,,: Transverse waves
Wave Propagation Simulations
- Helbold et al. 1986; Musielak, Rosner, Ulmschneider 1989,,, too many!
- Driving waves
✓
✓
Longitudinal piston-like motion
Transverse foot point motion
Wave Generation & Propagation Simulations
- Steiner 1998, 1999: Swaying; Bending; Squeezing
- YK et al. 2011: Magnetic Pumping
- Wedemeyer-Böhm et al. 2012: Swirling
ALMA太陽観測WS 2012/09/04
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5
s
g
n
i
l
e
d
o
M
c
i
t
Realis
Issues of Analytical Solutions
Thin flux tube approximation
- Wave equations become much simpler forms that can be solved analytically,
- Most of papers consider pores and magnetic elements as a flux tube, but this
approximation is valid only below the interface between convection zone and
photosphere,
- Some papers apply their results to sunspots instead of pores and magnetic
elements, but I think it s misleading.
Isothermal/Adiabatic atmosphere
- Radiation loss cannot be neglected in the Solar atmosphere, especially in the
interface between convection zone and photosphere, and also in the lower
chromosphere.
ALMA太陽観測WS 2012/09/04
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6
1982SoPh.
Issues of Wave Propagation
Simulations
Source of Waves?
- Location of wave generation remains mystery.
- Steady or Impulsive?
Wave Modes?
- Generation of waves is poorly understood in both
theory and observations.
- This is why there is a lot of choices and therefore
any drivers can be applied.
Energy Spectrum?
- Likewise, energy spectrum of generated wave is
unknown.
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What kind of MHD waves?
(Spruit 1982)
Ca II H-line formation
in 1-D Non-LTE Model atmosphere
by Carlsson & Stein 1992, 1994, 1995, 1997; Fossum & Carlsson 2006; Carlsson 2007
ALMA太陽観測WS 2012/09/04
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Ca II H-line formation
in 1-D Non-LTE Model atmosphere
by Carlsson & Stein 1992, 1994, 1995, 1997; Fossum & Carlsson 2006; Carlsson 2007
The asymmetry of the CaII H-line (H2V bright grains) is caused by high opacity and small overlying
opacity at the2012/09/04
H2V wavelength at the location of the shock.
ALMA太陽観測WS
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Propagation of Transverse Waves
Vigeesh, Hasan, Steiner 2009
G. Vigeesh
et al.: Wave propagation and energy transport in the magnetic network
955
Strong
Magnetic
Field
1200
500.0
1200
400
X (km)
800
200.0
1200
1
0.0
400
X (km)
400
1
1
0.0
Z (km)
1200
500.0
1200
400
1
1
400.0
0.0
200.0
800
1
1
-200.0
400
-250.0
-100.0
-500.0
400
800
β>1
a) V
X (km)
1200
0
0
-200.0
400
s
ALMA太陽観測WS
are excited at z = 0, where β > 12012/09/04
(on the axis β = 1.8), in the
from the base of the sheet cross this layer at some height. Waves
form of a fast (predominantly acoustic) wave and a slow (predominantly magnetic)1 wave, which propagate respectively at
Tuesday, September 4, 12
the sound and the Alfvén speeds. On the sheet axis, the acoustic
X (km)
1
1
200.0
0.0
150.0
800
1
1
800
b) Vn
1200
Time
β<1
0.0
400
-200.0
-150.0
1
1
300.0
1
1
1
X (km)
1200
Vn (m s-1)
0.0
1
-400.0
800
1200
Z (km)
1
400.0
400
-400.0
0
0
-300.0
1
400
1
β<1
1
0.0
800
Z (km)
1
100.0
0
0
-500.0
Vs (m s-1)
200.0
1200
0
0
400
800
1200
Fig. 3. Velocity
components
for the case
X (km)
in which the field strength at the axis
at z = 0 is 800 G (moderate field).
The 1200
colors (gray shades for the print
version) show the velocity components
a) V s , along the field, and b) Vn , normal
to the800field, at 40, 60, and 80 s (from
bottom to top) after initiation of an impulsive horizontal motion at the z = 0
boundary
of a duration of 12 s with an
400
amplitude of 750 m s−1 and a period of
P = 24 s. The thin black curves are field
lines and
0 the white curve represents the
800
1200
contour 0of β = 1.400TheXfield
aligned and
(km)
normal components of velocity are not
a) V s B < 50 G.
shown in the regions where
Z (km)
1200
250.0
Vs (m s-1)
Z (km)
1
X (km)
800
Vn (m s-1)
β<1
-200.0
400
1
1200
800
0
0
0.0
400
-250.0
1
X (km)
800
500.0
1200
1200
1
0
0
-500.0
400
800
1200
400
-100.0
1
0
0
X (km)
250.0
400
-250.0
-200.0
1
0
0
-600.0
800
800
800
-100.0
β>1
100.0
Z (km)
1
Vs (m s-1)
250.0
800
0
0
1200
Vn (m s-1)
X (km)
800
0.0
1
400
1
400
β>1
-300.0
1
-300.0
1
1
0
0
-700.0
0.0
800
Z (km)
-150.0
1
0
0
1
100.0
400
400
-350.0
Z (km)
0.0
1
Vs (m s-1)
1
1
400
1
Z (km)
0.0
800
800
Z (km)
1
Vn (m s-1)
1
200.0
1200
300.0
150.0
Z (km)
Z (km)
800
Vs (m s-1)
350.0
600.0
1200
Vn (m s-1)
300.0
1200
Vs (m s-1)
700.0
1200
Vn (m s-1)
et al.: Wave propagation and energy transport in the magnetic network
WeakG. Vigeesh
Magnetic
Field
400
wave vector and the local direction
of thestratification,
magnetic field
gravitational
the(Cally
β = 1 surface (more precisely the
2007). On the β = 1 layer, away
from
the
sheet
axis,
where
thespeed) would still constitute
surface
of equal Alfvén and sound
9
wave vector is not exactly parallel
to the
magnetic
field,
we do so that we could expect scethe critical
layer
for mode
coupling
not have complete transmission
of the
wave different
to a slowfrom
wave.but more complex than those
narios
not fast
radically
Rather, there is a partial conversion of the mode from fast acous-
X (km)
800
b) Vn
1200
Fig. 5. Velocity co
in which the field
z = 0 is 1600 G (
t = 40, 60, and 8
sponds to that of
(from bottom to top) for the moderate fi
ambient medium the field strength is wea
are not shown in this region. The Poynt
the wave energy that is carried by the ma
Time Evolution of Flux Tube in the
Solar Atmosphere
YK, Steiner, Steffen, Suematsu (2011) by using CO5BOLD
white contour: Magnetic field line, Arrows: velocity
White & Black curve: Optical surface
Tgas
Nx=400 Nz=165
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EXCITATION OF SLOW-MODE & SHOCK
YK, Steiner, Steffen, Suematsu (2011)
Downflow
-10
-5
0
Vz [km s-1]
5
Upflow
10
time = 12.5 min
time = 13.8 min
time = 15.1 min
time = 16.4 min
200
0
z [km]
-200
core
-400
-600
-800
-1000
-1200
4000 4400 4800 5200 4000 4400 4800 5200 4000 4400 4800 5200 4000 4400 4800 5200
(a)
x [km]
x [km]
x [km]
x [km]
z [km]
nflo
flo
w
up
aft
dr
wn
skin
do
skin
dow
0
-200
-400
-600
<τ>core = 1
<β>core = 1
-800
-1000
(b)
ALMA太陽観測WS 2012/09/04
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-1200
12
11
shock
<Vz>core [km s-1]
w
<Vz>skin [km s-1]
200
Spatially averaged vertical
velocity in the skin region
13
14 15 16
Time [min]
Spatially averaged vertical
velocity in the core region
17 18 12 13
14 15 16
Time [min]
17
18
EXCITATION OF SLOW-MODE & SHOCK
YK, Steiner, Steffen, Suematsu (2011)
Downflow
-10
-5
0
Vz [km s-1]
5
Upflow
10
time = 12.5 min
time = 13.8 min
time = 15.1 min
time = 16.4 min
200
0
z [km]
-200
-400
-600
-800
Shock
wave
-1000
-1200
4000 4400 4800 5200 4000 4400 4800 5200 4000 4400 4800 5200 4000 4400 4800 5200
(a)
x [km]
x [km]
x [km]
x [km]
z [km]
-400
-600
(b)
ALMA太陽観測WS 2012/09/04
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-1200
12
11
shock
flo
w
<τ>core = 1
Magnetic pumping
<β>core = 1
-800
-1000
w
up
aft
dr
wn
do
-200
Downflows
dow
Downdrafts
0
<Vz>core [km s-1]
nflo
<Vz>skin [km s-1]
200
Spatially averaged vertical
velocity in the skin region
13
14 15 16
Time [min]
Spatially averaged vertical
velocity in the core region
17 18 12 13
14 15 16
Time [min]
17
18
CORRELATION BEWTEEN
DOWNDRAFTS AND SLOW-MODES/SHOCKS
YK, Steiner, Steffen, Suematsu (2011)
Formation of shock at
Chromospheric height
Propagation speed 5 - 10 km/sec
Downflows are excited near
the optical surface (tau 1)
-10
-5
0
5
10
<Vz>core , <Vz>skin [km s-1]
z [km]
500
0
-500
<τ>core = 1
-1000
0
20
40
Time [min]
Downdrafts are excited below
the optical surface (tau 1)
ALMA太陽観測WS 2012/09/04
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60
80
磁束管で自発的に生じる
磁気流体波とは？
Longitudinal
wave
Corona
Chromosphere
β=1
Photosphere
τ=1
Convection
zone
ALMA太陽観測WS 2012/09/04
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Transverse
wave
Torsional
wave
Pumping Buffeting Swirling
13
Magnetic Tornadoes
in the Solar Atmosphere
Wedemeyer-Böhm et al. (2012) by using CO5BOLD
ALMA太陽観測WS 2012/09/04
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What we can observe by using
ALMA?
ALMA: 4000 - 7500 K
ΔX=28km
Nx=400 Nz=535
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ALMAによる多層間結合
高解像高速撮像
[email protected]
ALMA: 4000 - 7500 K
- Higher-spatial resolution
10 km
- Internal structure of the base of flux tube (DC)
- Propagation of shocks along the flux tube (Limb)
[email protected]
- Comparable spatial resolution
30 km to models
- Detect propagation of shocks inside the flux tube
[email protected]
- Lower-spatial resolution
100 km
- Shock propagation time 1000 km / 10 km/s = 100 s
- High-cadence observation is useful < 1 min
ALMA太陽観測WS 2012/09/04
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まとめ
動的な太陽大気モデリングにはRMHDが必要不可欠！
磁気大気：Canopyなどの複雑な磁場構造によって大気構造が決まっている
磁束管で自発的に遅い磁気音波が発生
上部対流層での速い下降流(downdraft)によって遅い磁気音波が励起
光球から上空へ遅い磁気音波が伝播し、彩層で衝撃波となる
3次元ならは渦を伴うトルネード状になるだろう
磁束管で光球・彩層の振動現象の存在を予言
高解像度(< 0.1 )高頻度(< 60sec)の光球・彩層観測が必要
磁束管の構造を知るには、光球・彩層磁場観測が必要不可欠
Solar-Cの分光偏光観測とALMAの高解像度撮像に期待
ALMA太陽観測WS 2012/09/04
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