UCL DEPARTMENT OF SPACE AND CLIMATE PHYSICS
MULLARD SPACE SCIENCE LABORATORY
Energy release and transport in solar
flares & CMEs
Sarah Matthews
Red RHESSI 6-12 keV, blue 50-100 keV, gold images TRACE 195A
Flare - definition and classification
A solar flare is a sudden release of energy during which via magnetic reconnection free
magnetic energy is converted to kinetic energy of fast particles, mass motions, and
radiation across the entire electromagnetic spectrum.
Energy released up to 1025 J or 1032 erg in the largest solar flares.
Many more much smaller flare-like events occur, with energies of as small as 1016 J – nanoflares, micro-flares etc.
GOES soft X-ray classification is most common
these days.
The flux in the 1–8 Å = 0.1–0.8 nm range is
recorded by this scheme:
B → 10-7 W/m2
C → 10-6 W/m2
M → 10-5 W/m2
X → 10-4 W/m2
-  so an M5 flare has flux of 5 × 10-5 W/m2.
GOES = Geostationary Environmental
Operational Satellites – they are continuously
recording solar X-ray emission. Flare time profile
Impulsive phase – primary energy release
•  hard X-rays (10s of keV)
•  white light, UV, µwaves - broad spectrum
•  duration < few minutes
•  intermittent and bursty time profile, 100ms
•  energy injection:
- few tenths of the total flare energy released
(up to 1032 ergs)
- significant role for non-thermal electrons
Gradual phase - response to input
•  thermal emission (kT ~0.1-1 keV)
•  rise time ~ minutes
Neupert effect t(Neupert, 1968):
FSXR (t) ≈
∫F
HXR
(t ' )dt '
t0
€
Soft X-rays mainly originate from plasma heated
by the accumulated energy deposited by
accelerated electrons from flare start.
Flare frequency spectrum
The number of flares falls off with
increasing power as a flat power law
with a slope of ~ -1.8
(SXR, EUV, microwave, HXR bursts,
optical flares) (e.g. Drake, 1971; Dennis,
1985, Hudson, 1991)
dN/dW=A·W-α (ergs s)-1
<1-hour dataset!
the normalisation factor A varies with
the level of activity (Kreplin et al, 1977,
Wagner, 1988)
α=-1.8
Recently α values were found
between 1.5 and 2.6.
Relevance for coronal heating:
Aschwanden et al, (2000)
If α<2, smaller flares do not contribute
enough to heat the corona. Light curve of typical superflares.
H Maehara et al. Nature 000, 1-4 (2012) doi:10.1038/nature11063
Comparison between the occurrence frequency of superflares on G-type stars and those of solar
flares.
Kazunari Shibata et al. Publ Astron Soc Jpn 2013;65:49
© Astronomical Society of Japan
Magnetic indicators of imminent flares?
•  Intuition:
–  Complex, rapidly evolving active regions have the highest
probability of flaring
•  More quantitative methods?
–  Past X-ray activity (Bayesian stats)
•  Wheatland (2004) – moderately successful
–  Magnetic field stats and variations
•  No consistent picture emerged (Leka & Barnes, 2006)
•  >70% data flare quiet at C1.0 level, so by doing nothing you get >70%
success rate….
Prior flaring
Falconer et al. (2012)
Red – prior flaring; Blue – no prior flaring
High-gradient + strong-field - flaring
Schrijver (2007) introduced R (new metric): summed unsigned B of high-gradient strong-field polarity
inversion line. (overlap of +ve or -ve B > 150 Mx cm-2 Kernels: 6”x6”
SOHO/MDI Φ= R x 2.2x1016 Mx
Schrijver, 2007 Forecast of major flare within 24 hours:
R ≥ 2x1021 Mx (logR≥4.8), probability ≈ 1
R ≤ 1019 Mx (logR≤2.8), probability ≈ 0
These features are characteristics of new flux
emergence in highly non-potential state. ◊: M&X flares Strong R, Φ & big flare correla2on Origin and storage of free magnetic energy
Origin:
•  Magnetic flux emerges twisted, i.e. in a non-potential state, from the solar interior.
Twist may keep propagating from below via torsional Alfvén waves.
•  Surface flows and magnetic footpoint shuffling shear and entangle field lines.
Storage:
•  Magnetic free energy (above the energy of the potential state) is stored relatively
low in an AR ≤ 20 Mm above the photosphere and may mainly be concentrated
along the magnetic inversion line in the filament channel.
Hinode/SOT magnetogram
Integrated electric currents in an AR before (a) and after
an X-class flare. Note their organization into an apparent
flux-rope structure (Schrijver et al., 2008)
Twisted flux emergence and X-class flare
Hinode/SOT
G-band
Upper photosphere
WLF < 100 km height
Isobe et al, 2007
Hinode/SOT
Ca II H-line
Chromosphere
Kubo et al., 2007
How is the energy released?
•  Magnetic reconnection is a topological restructuring of a magnetic field caused
by change in the connectivity of its field lines.
•  It allows the release of stored magnetic energy (dominant free energy in plasma)
• Evidence of reconnection has now (we believe) been seen
Observation of the energy release site remains controversial
Many pieces of indirect evidence in solar observations.
CSHKP model for eruptive solar flares
(Carmichael 1964; Sturrock 1966; Hirayama 1974; Kopp & Pneuman 1976)
Filament post reconnection
prereconnec
tion
prereconnec
tion
Adapted from Shibata (1998)
post reconnection
•  Accelerated electons gyrate
along magnetic field lines
emitting gyrosynchrotron
radiation
•  Collisions in the dense
chromosphere emits
bremstrahlung observed in
hard X-rays (> 20 keV).
•  Electrons impulsively heat the
chromosphere leading to
optical and UV emission.
•  Heated chromospheric
plasma expands upward,
increasing ρ and T in the
reconnected coronal loops.
Confined flares - quadrupolar reconnection
Reconnection happens
•  at nullpoints (X-point)
•  at separatrices and their intersection, the separator
•  at quasi-separatrix layers (QSLs)
Signatures:
•  four flare kernels/ribbons at the footpoints of
reconnected loops
•  ribbons are in the vicinity of drastic field line
connectivity changes.
Priest & Forbes, 2000
Reconnection along QSL
•  Along QSLs field line mapping is
continuous but shows steep
gradients.
•  Reconnection along QSL does
NOT break and reconnect field
lines, but field lines may slip
across each other, as shown in
MHD simulations.
•  The movie shows a case for sliprunning reconnection observed
with Hinode/XRT.
(Aulanier et al., 2007)
Su et al., 2013
Rising reconnection region, cooling loops
Hinode/EIS: Cooler loops lie below the hotter loops since the lower ones
were formed before the higher ones
Shrinking reconnected loops
First examples were found in
Yohkoh/SXT data
McKenzie & Hudson (1999)
TRACE EUV
•  Shrinkage of newly reconnected cusped loops driven by the magnetic
tension force.
•  Patchy and intermittent reconnection process.
Patchy and intermittent reconnection
Start of these downflows are
associated with non-thermal HXR
emission (RHESSI) and microwave
bursts (NoRH). Asai et al., 2004 Su et al.,
2013
In 3-D
Musset et al., 2015
Janvier et al. 2014
Zharkov et al., 2011
How & where are particles accelerated?
•  New insights from 3D
models:
–  Coronal X-ray emission overlies
current ribbons in the
photosphere
–  New >50 keV HXR source
appears in association with
increased photospheric current
–  > clear link between particle
acceleration and reconnecting
current sheets
Musset et al., 2015
Hard X-rays
•  Produced by electron-proton bremsstrahlung from electrons
>15 keV
1
I (ε) =
nV
2
4 πR
∞
∫ F (E)Q(ε, E )dE
ε
•  Thermal bremsstrahlung: Eelectron ~ Etarget and spectrum F(ε) ~
e-ε/kT
•  Non-thermal
bremsstrahlung: Eelectron >> Etarget and spectrum
€
F(ε) ~ ε-γ
•  Very inefficient, ~ 10-5 electron energy radiated as X-rays
•  In impulsive phase, HXR spectrum can be fitted by a hot (20
MK) or superhot (~60 MK) thermal component plus a power
law.
time
frequency
frequency
Isliker & Benz 1994
Radio emission
time
Upward and downward-going beams are observed, occurring at peak time of
HXR emission.
Metric and decimetric Type III bursts are often plasma radiation produced by
electron beams (from Langmuir waves at f ~ 9 √ne).
Impact of accelerated particles
Precipitation of energetic electrons
and ions from the coronal
acceleration site in the dense
chromosphere →
heating → overpressure →
plasma upflows (chromospheric
evaporation)+ bright flare ribbons +
HXR & γ-ray sources along the
ribbons.
Gamma-ray footpoints
•  HXRs and gamma-ray lines have similar time
profiles, implying related acceleration, but ion
signature is in different location from electron
signature.
Neutrons produced by energeCc ions (10s of MeV/nucleon. Capture line predicted to form within 500 km of neutron producCon site. Observed offset from HXRs, ~10000 km Hurford et al., 2006 Flare quakes
•  Seismic disturbances seen
during the impulsive phase
of a small number of flares.
•  First detected by Kosovichev
& Zharkova (1996)
•  Approximately co-spatial
with HXR and WL emission.
•  How does the energy get so
deep? 15 scale heights!
Kosovichev & Zharkova, 1998
Kosovichev & Zharkova, 1998
Flare related magnetic field changes
•  Magnetic reversals are
seen in some flares,
spatially and temporally
correlated with HXR
sources – flare related
changes to the line
profile?
•  Also sudden and
‘permanent’ changes in
the longitudinal field of
~ 10% (100-200 G)
(Sudol & Harvey (2005).
•  Some are co-spatial
with flare ribbons/
kernels and propagate
at flare ribbon speed.
Johnstone et al., 2009
Magnetic re-structuring and Alfvén waves
•  Alfven waves heat
chromosphere & drive
evaporation (Reep &
Russell, 2016)
!  Increased line widths,
coronal upflows, lack of
HXR signature
•  Restructuring causes
increase in horizontal B/
change in tilt
•  Increased photospheric
currents
Russell et al., 2016
Reep & Russell, 2016
Evidence?
•  Broadening of chromospheric lines/
absence of co-spatial HXR
emission
–  vD =134.7 km s-1 => wave
energy flux ~ 1011 erg cm-2s-1
•  Increased j at SQ location
Mg II k
Matthews et al., 2015
Sharykin et al., 2015
CMEs
•  “observable change in coronal structure that occurs on
a timescale of a few minutes to several hours, and
involves the appearance and outward motion of a new,
discrete, bright, white light feature in the coronagraph
field of view.” (Hundhausen, 1986)
First observations
Eclipse drawing (Tempel) 18 Jul 1860
Skylab 10 June 1973
(MacQueen et al. 1974)
Eddy 1974
Morphology
•  3 part structure:
–  Bright fontal
loop (overlying
arcade)
–  Dark cavity (flux
rope)
–  Bright core
(filament/
prominence)
•  ~30% show this
structure
Jets, Halos and
partial halos
•  Narrow jet-like
structures ~20°
•  Partial and full
halos (120
-360°)
•  Directed along
the Sun-Earth
line.
•  10% of all CMEs
are halos; 4%
full.
Thomson scattering
• CMEs are best observed in
WL – Thomson scattering of
photospheric light by coronal
electrons
• Depends on density of
scattering electrons and angle
between incident radiation
direction and the l.o.s.
• Scattering is strongest in the
plane of the sky, i.e. limb
CMEs are favoured.
Vourlidas & Howard, 2006
Different perspectives
•  Single viewing perspectives can be misleading and
give skewed perception of angular widths and other
properties.
•  70° separation of STEREO A and B from LASCO.
STEREO A
LASCO
STEREO B
Radio signatures
Bastian et al., 2001
CME speeds > vA drive a shock
ahead of the CME which can accelerate
electrons -> Langmuir waves -> Type II radio bursts
Properties
• 
• 
• 
• 
Contain coronal material ~ few MK
Cool prominence material in the core ~ 8000 K
Compressed sheath behind the shock higher T and ne.
Cavity lower density -> higher magnetic field for
pressure balance. Virtually no measurements…
Angular width
•  Extent of PA in plane
of the sky => only get
accurate results for
limb CMEs.
•  Widths range from 5 360°.
•  Average computed for
< 120°.
Speed and Acceleration
•  CMEs start from rest.
•  Driving force close to
the Sun, interaction
with solar wind slows
the CME.
•  Speeds measured
from fits to H-t plots
•  Fast CMEs
decelerate, slow
accelerate.
•  a = -0.015(V-466)
LASCO 1996 -2009 (Gopalswamy 2010)
Kinematics
•  Mass: estimate number of electrons needed in plane of
the sky to produce observed brightness.
•  M = (Bobs/Be(θ))x1.97x10-24 g; Be(θ) – brightness of a
single electron at angle θ from the plane of the sky
(Vourlidas et al., 2000).
•  Speed – fits to height-time measurements in the plane
of the sky (more reliable for limb than halo)
•  V =360 +3.64 W (W – angular width)(Gopalswamy et
al. (2009))
•  Wider CMEs are generally faster and more massive.
Mass and KE
Figure 6. Distributions of CME acceleration obtained from a quadratic fit to the height – time measurements for
(left) and limb CMEs (right). The period considered: 1996 – 2009.
• Faster and wider
CMEs have higher KE
• KE ranges from < 1026
to > 1033
• Distributions become
more symmetric if only
limb CMEs are
considered, and median
values increase:
7. Mass
30and kinetic energy distributions of SOHO/LASCO CMEs (1996 – July 2008). The median val
1.3x1015g and Figure
1.6x10
distributions
are marked on the plots.
Gopalswamy 2010
ergs.
ourselves to limb CMEs, as we did when estim
2.4.4 Mass and Kinetic Energy
The CME mass (computed as the excess mass in the
coronagraphic FOV) ranges from <1012 to >1016 g with
a median value of ~3.2x1014 g (see Fig. 7). The CME
mass is computed by estimating the number of electrons
needed in the sky plane to produce the observed CME
CME acceleration, the distributions of mass an
energies become more symmetric and the medi
become several times greater: 1.3x1015 g and
erg, respectively. The higher values are consi
pre-SOHO values (see e.g., Howard et a
because those coronagraphs were not sensitiv
Statistical properties
Table from Webb and Howard, 2012
Flares and CMEs
•  Both arise due to reconfiguration and subsequent
release of energy in the coronal magnetic field.
•  Plenty of flares which are not accompanied by CMEs,
or are failed eruptions.
•  CMEs can also occur without flares, although many
models predict flares as part of the eruption process.
•  Correlation increases with flare size. 100% for >X3.0
(Yashiro et al., 2005)
On-disk signatures
•  Flare ribbons and
arcades
•  Coronal dimming
•  EUV waves
Patsourakos et al.,
2009
Occurrence rate
•  Rate varies from
<0.5/day and
minimum to > 6 at
max.
•  Can be > 10/day
•  Significant
differences
between
catalogues –
automatic vs
manual; observer
bias; sensitivity.
Robbrecht et al., 2009
Solar cycle variation
• CME occurrence rate
follows the solar cycle in
phase and amplitude
• Linear relationship
between CME rate and
sunspot number (Webb &
Howard, 1994; Robbrecht
et al., 2009)
• But correlation is higher
during rise and declining
phases than during max.
• Speeds are generally
higher at max.
Robbrecht et al. 2009
Latitudinal distribution
Gopalswamy 2010
Halos
Halo speeds
Shocks
• Fast CMEs drive shocks
• Shock accelerates
electrons, which produce
Langmuir waves.
• Converted to Type II radio
emission at plasma
frequency.
• Drifts to decreasing
frequency -> density.
• Direct imaging of shock
sheath only really just been
recognized – cloud of
electrons ahead of the
CME.
Rouillard et al., 2012
CMEs and SEPs
• Gradual SEP events
strongly associated
with CMEs
• Impulsive – flares
• Also hybrid events
• Sources in West
magnetically well
connected to Earth
• But CME flanks are
wide.
Reames (1999)
CMEs and SEPs
•  STEREO is
showing that SEP
events have up to
360° impact.
Dresing et al. 2012
EUV waves
Carley et al. 2013
Interplanetary CMEs
Outflows and magnetic
clouds
•  Spectroscopic
measurements of outflow
velocity have been used
to compute more reliable
estimates of magnetic
flux in a CME
•  Good agreement with
estimates of flux in
associated magnetic
cloud.
Harra et al., 2011
Models
After eruption onset:
•  Observations of flux ropes
• 
• 
• 
at 1AU,
CME acceleration during
flare impulsive phase,
flare ribbons and flare
arcade loops,
EUV & X-ray dimmings.
The above has led to the
‘standard model’
Figure 1 from Shibata et al. (1995)
Mandrini et al., 2005
In 3-D
Musset et al., 2015
Janvier et al. 2014
Zharkov et al., 2011
Aly-Sturrock
•  If all field lines simply linked to the solar surface, total
energy of any force-free field cannot exceed that of an
open field with the same flux distribution on the
surface.
•  Implies ‘opening’ the field energetically unfeasible.
•  Possible solutions:
–  Reconnection
–  Initial configuration not simply linked to the surface
–  Initial field not force-free – currents, gravity
–  Only part of the field is opened
–  3D - flux rope can erupt by pushing the field aside.
Triggering and models
•  Energy must come from the magnetic field
•  Pre-CME equilibrium maintained by the balance of
magnetic tension and magnetic pressure.
•  Decrease of tension or increase of pressure could
cause the pre-CME structure to seek a higher
equilibrium and a CME may be triggered.
•  Many different mechanisms proposed:
–  Resistive
–  Ideal
Resistive Models: Tether cutting
• Filament supported by field
nearly aligned with the PIL
• Strongly sheared ‘core’ field
near the PIL
• Overlying less sheared arcade
• As shear increases
reconnection starts.
• Tethers (AB and CD) are cut
and AD and CB are formed.
• AD expand upwards and CB
submerges, while sheared field
near AD pulls the filament up.
• Overlying field is stretched,
forming a current sheet.
• Doesn’t address formation of
the sheared field.
D
A
C
B
Moore et al., 2001
Resistive models: flux cancellation
By emergence:
By reconnection:
B.C. Low, 1996
van Ballegooijen & Martens, 1989
Debated!
Currently favoured:
accompanied by photospheric flux
cancellation
How does the flux rope form?
• 
• 
• 
Successive reconnections along the PIL (van Ballegooijen &
Martens, 1989)
Is BPSS or HFT formed?
Recent simulation of the process by Aulanier et al. (2010)
Resistive Models: Break-out
• Quadrupolar configuration
with a null point above the
central flux system.
• Shearing motions cause the
central flux system to expand
upwards, forming a current
sheet at the null point.
• Reconnection starts,
removing higher magnetic
loops and allowing the core
field to erupt.
• Kind of external tether cutting.
Antiochos et al., 1999
Ideal models: Kink instability
•  Can explain height-time profile of an erupting filament
(Sakurai (1976))
•  Kink: critical twist 2π to 6π (Hood & Priest, 1979)
•  Török & Kliem (2005) considered overlying field:
–  Rapid decay with height leads to eruption
–  Slow decay with height – failed eruption
Fan & Gibson (2007)
Ideal models: Torus instability
•  Current ring is unstable against expansion if external
field decay is fast.
•  Hoop force dominates over magnetic tension and flux
rope can no longer be confined (Kliem & Török, 2005)
Fan & Gibson (2007)
Summary
•  Flares and CMEs represent the largest release of
• 
• 
• 
• 
magnetic energy within the solar system.
Proximity of the Sun and many observing platforms
allow detailed testing of ideas about how energy is
released and transported.
Knowledge can be used to inform understanding of
other astrophysical environments.
The ‘standard’ model has many successful elements,
and a 3-D picture is now well developed.
Observational improvements show inconsistencies –
the problem isn’t solved yet.