particle acceleration in solar flares based on hinode flare

Particle Acceleration in Solar Flares based on Hinode Flare Observations

Kyoko Watanabe

ISAS/JAXA

Hinode6 @ St. Andrews　 　 　2012 Aug 14

Particle Acceleration in Solar Flares based on Hinode Flare Observations

Kyoko Watanabe (ISAS/JAXA)

1.  Introduction 2.  Multiple wavelength observations of Solar Flares

by Hinode, hard X-ray and radio 3.  White-light flare （by SOT & RHESSI） 4.  Electron injection to the foot-points and Plasma

diagnosis（by EIS & RHESSI） 5.  Particle acceleration from microflares 6.  Comparison of ion and electron acceleration and

solar neutron observations 7.  Future studies of particle acceleration by Hinode

Hinode6 @ St. Andrews　 　 　2012 Aug 14

magnetic reconnection model

1. Introduction: Particle acceleration in solar flares

Total energy of solar flare: 1029 - 1033 erg (1022 - 1026 J)

•  plasma heating •  electromagnetic radiation •  particle acceleration - Electron: ~ hundred MeV - Ion: ~ GeV

•  Electron acceleration - Radio - hard Ｘ-ray - Bremsstrahlung γ-ray - White-light

•  Ion acceleration - Nuclear γ-rays - Solar neutron

•  Electron acceleration - Radio: NoRH etc.. - Ｘ-ray, bremss.γ-ray: RHESSI, Fermi etc. - White-light: TRACE, Hinode, SDO, Ground Telescopes

•  Ion acceleration - Nuclear γ-rays: RHESSI (<10MeV), Fermi (π0) - Solar neutron: Neutron Monitor, SEDA-AP (ISS)

1. Introduction: Particle acceleration in solar flares

•  Most of them have the capability of observing the full Sun •  They keep on observing except during night time and SAA •  Hinode has a limited FOV for observing with high resolution and high cadence •  To know when and which solar flares are observed by Hinode’s instruments is very useful and helpful for scientists in surveying flares for analysis

Solar flare observations by Hinode （Watanabe et al., 2012）

Hinode Flare Catalog

EIS

XRT

SOT

328

164

1024

590

2048

2048

http://st4a.stelab.nagoya-u.ac.jp/hinode_flare/

Not all flares are seen by Hinode

Flare occurs when Hinode is observing + Inside Hinode FOV

⇩

1. Introduction: Particle acceleration in solar flares

XRT SOT (FG) EIS Total event # 2006 (10/25~) 160 171 68 432

2007 443 267 209 713 2008 24 8 7 134 2009 125 83 65 320 2010 690 264 125 1248 2011 1269 464 276 2389

total 2711 (51.8%)

1257 (24.0%)

750 (14.3%)

5236

GOES class XRT SOT (FG) EIS Total event # X 7 5 3 12 M 81 29 28 144

Number of solar flares observed by Hinode 1. Introduction: Particle acceleration in solar flares

2.  Multiple wavelength observations of Solar Flare by Hinode, hard X-ray and radio

•  Determined the location of particle acceleration in the corona •  Examined the pitch-angle distribution of accelerated electrons

X-class flare on 13 Dec 2006

Hinode flare =

-  Minoshima+ 2009 Hinode/SOT + Hinode/XRT + RHESSI + NoRH + NoRP

Minoshima et al., 2009 Contours: RHESSI 35-100 keV at 02:29:04 UT Dotted lines: Ca II H ribbons at 02:30:38 UT

Bx By Bz

HXR sources are located at the regions (1) with the strongest magnetic field along the ribbons (2) only where the direction of the horizontal magnetic fields change remarkably.

Position of the hard X-ray sources ⇒ position of particle acceleration in the Corona

Grey-scale: XRT Grey contours: NoRH 34GHz Black contours: RHESSI 35-100 keV

They interpret the spatial distribution of the microwave intensity as resulting from electron injection parallel rather than perpendicular to the magnetic field lines.

Pitch angle distribution of accelerated electrons

Minoshima et al., 2009

- Krucker et al., 2010 RHESSI + NoRH + NoRP + STEREO + Messenger + SOHO/EIT + GOES

2.  Multiple wavelength observations of Solar Flare by Hinode, hard X-ray and radio

Multiple wavelength & multiple viewpoint observations of high energy coronal source

The hard X-ray and microwave emission in the impulsive phase occurs at higher altitude than the thermal loops seen during the soft X-ray peak time

EIT: 195A RHESSI: 6-8keV RHESSI: 30-50keV NoRH: 17GHz

Krucker et al., 2010

Impulsive phase Soft X-ray peak

1.  The acceleration site is in the coronal low-β plasma located above (~6000 km) the main thermal flare loops seen in soft X-rays.

2.  The acceleration region does not emit significantly at EUV and soft X-ray wavelengths.

3.  All of the electrons are accelerated into a power-law distribution extending up to the MeV range.

Krucker et al., 2010

Hudson et al. 2006

Fletcher et al. 2007

White-Light emissions are well correlated with hard X-ray and

radio emissions (Location & Profile) → Non-thermal electrons

3. White-light flare

White-light flare observations: •  TRACE（WL+UV） + RHESSI •  SOT (G-band, RGB) + RHESSI •  SDO/AIA, HMI + RHESSI

Very rare event observed only from large solar flares (X-class etc.)

(Hiei 1982, Neidig 1989) ↓

White-light emissions from C-class flares also observed thanks to

accurate photometry by Satellites (Matthews et al. 2003, Hudson et al. 2006)

White-light flares observed by Hinode/SOT （X-class）

(Wang 2006, etc.)

Date GOES class

flare location

SOT

2006/12/06 X6.5 S05E57 G-band 2006/12/13 X3.4 S07W22 G-band 2006/12/14 X1.5 S06W46 G-band 2011/02/15 X2.2 S20W10 R, G, B 2011/11/03 X1.9 N21E64 R, G, B 2012/01/27 X1.7 N33W85 R, G, B 2012/03/05 X1.1 N19E58 R, G, B

White-light emissions from M- and C-class also observed by Hinode.

2006 Dec 6 White-light flare SOT/G-band + RHESSI (Krucker et al., 2011)

The hard X-ray and white-light (G-band) data show very similar ribbon structures (length ~30”), and this result strongly

suggests that the flare emissions in white light and in hard X-rays have physically linked emission mechanisms.

2006 Dec 14 White-light flare SOT/G-band + RHESSI (Watanabe et al., 2010)

White-light & hard X-ray emissions were seen at almost the same location.

•  Power of accelerated electrons 　→ Thick Target Model •  Power of White-Light emission 　→ Blackbody

The power of white-light (G-band) emission can be explained by >40keV non-thermal electrons.

Other results of previous studies: >50keV: Neidig (1989) etc. >25keV: Fletcher et al. (2007)

Difference of the emission height of WL & HXR

0

600

km 6500

Chro

mos

pher

e Co

rona

Phot

osph

ere

G-band (Carlsson et al., 2007) Continuum

50-100keV HXR (From Yohkoh events, Matsushita et al., 1992)

50-100keV HXR (by RHESSI event, Kontar et al., 2008)

100

1000 Thermalize 50-100keV electrons (density = 1013.5/cm3, Neidig 1989)

The data suggest a difference of more than 500km between the emission sites.

>900keV are needed to reach the photosphere (Neidig, 1989)

↑6173 Å continuum (AIA) @ 1.5 – 3 Mm (Battaglia & Kontar, 2011) Ca H

G-band

2011 Feb 24 White-light flare SDO/AIA, HMI + RHESSI (Battaglia & Kontar, 2011)

WL (black) emission located higher than HXR (blue,

25-50keV) emission ↓

Origin of WL emission is low energy electrons (<12keV)

2011 Feb 24 White-light flare SDO/AIA, HMI + RHESSI (Oliveros et al., 2012)

Emission height - White-light: 195±70 km - Hard X-ray: 305±170 km

Orange contours: RHESSI 6-8 keV Blue contours: RHESSI 30-80 keV Magenta contours: SDO/HMI 6173.3Å

Accelerated e- (30-80keV) reach the very near photosphere.

Both upflows and downflows were spatially and temporally correlated with HXR emission. RHESSI was used to derive the properties of the electron beam deemed to be the driver of the evaporation. The energy flux of the e- beam ~ 5 × 1010 erg/cm2/s

4.  Electron injection to the foot-points and Plasma diagnosis（by EIS & RHESSI）

(Milligan & Dennis, 2009; Milligan 2011)

5.  Particle acceleration from microflares

Small energy release phenomena (Hinode observed): -  Transient bright jets -  X-ray bright points etc…

→ no evidence of particle acceleration

•  RHESSI results: hard X-ray microflares - Christe et al., 2008 Statistical analysis for 25,705 events - Hannah et al., 2008 (spectral fitting: 9161 events)

•  Radio bursts - Morioka et al., 2007 (micro-type-III) - Iwai et al., 2012 (type-I)

Smallest energy release phenomena ever observed?

Hard X-ray microflare observations by Balloon

Observation for 141 min >3σ enhancements: 25 events

Lin et al., 1984

Microflare occurrence rate Power-law distribution

(same as flare) Spectra: power-law (γ~ 4 – 6)

Christe et al., 2008

2002 Mar - 2007 Mar: 25,705 events

Average: GOES A-class

Total energy flux of microflare-associated accelerated electrons (>10 keV): <1026 erg/s

RHESSI microflare observations (frequency distribution)

•  No relationship between loop-length and flare-energy •  Steep hard X-ray spectrum (γ = 7) •  Frequency distribution cannot be fit by power-law (total 9161 flares)

RHESSI microflare observations (spectrum distribution)

Hannah et al., 2008

Morioka et al., 2007

High time resolution observation by Akebono

Micro-type-III bursts　 vs

Ordinary type-III bursts

Frequency dist.: power-low → different index

Micro-type-III bursts

Type-I bursts Iwai et al., 2012

Frequency distribution of 190MHz on 2010 Jan 26

The frequency distribution was also shown to be a power-law whose index was larger than 2. 　⇒ Solar flare: α<2

Type-I bursts are different from solar flares?

(Hurford et al., 2003; Lin et al., 2003) Observation of X4.8 flare on 2002 Jul 23

The first 2.2MeV γ-ray line image from RHESSI

20” difference

Difference of acceleration and /or propagation mechanisms

6.  Comparison of ion and electron acceleration

(Hurford et al., 2006)

γ-ray line observation on 2003 Oct 28(X17), 29(X10), Nov 3(X8)

Observed 2.2MeV source is very compact 　→ accelerated ions originate from solar flare (not CME) Location difference of 200-300keV and 2.2MeV emissions 　→ difference of acceleration and/or propagation mechanisms

2003 Oct 28 X17

2003 Oct 29 X10

6.  Comparison of ion and electron acceleration

(Shih et al., 2009)

Statistical analysis of 29 γ-ray line (2.2MeV) events (2002~2005)

>300keV e- vs >30MeV p → well correlated

Correlation with soft X-ray → weak correlation

6.  Comparison of ion and electron acceleration

6.  Comparison of ion and electron acceleration Fermi/LAT, GBM γ-ray observation (Ackermann et al., 2011)

6.  Comparison of ion and electron acceleration Fermi/LAT γ-ray observation (Tanaka et al., 2011; etc.)

Π0 γ-ray emission lasted for ~12 hours.

No evidence of long time γ-ray emission by RHESSI.

Solar neutron Observation: SEDA-AP SEDA-AP: Space Environment Data Acquisition equipment– Attached Payload (SEDA-AP)

SEDA-FIB

BBD detector (<30MeV)

FIB detector (30-120MeV)

2009.12.03 23:59:59UT background neutron Typical example of observed neutron signal

Solar neutron Observation: SEDA-AP

  We have searched for solar neutrons in association with all flares with an intensity higher than M2 during 2011.

  We have found some neutron signals from an M3.7 flare on March 7, 2011.

7.  Future studies of particle acceleration by Hinode

•  Collaborating with observations of non-thermal emissions (hard X-ray, radio etc.) = RHESSI, NoRH …

•  Research using Hinode capability 　 → Understanding the location and field condition of accelerated particles. -  Temperature analysis by XRT and EIS (temperature, density) -  3D coronal magnetic fields estimated from SOT photospheric magnetic field (field strength)

•  Future mission •  Focusing optics in hard X-ray provides much higher dynamic range and sensitivity (FOXSI, NGXT) •  Chromospheric magnetic field observation