001–008

#

(2009)

††volume: Volume XXVIIA††journal: Transactions IAU, Volume XXVIIA††editors:
Karel A. van der Hucht, ed.

COMMISSION 10 | SOLAR ACTIVITY
---|---
|  ACTIVITÉ SOLAIRE

PRESIDENT | James A. Klimchuk
---|---
VICE-PRESIDENT | Lidia van Driel-Gesztelyi
SECRETARY | Carolus J. Schrijver
PAST PRESIDENT | Donald B. Melrose
ORGANIZING COMMITTEE | Lyndsay Fletcher, Nat Gopalswamy,
| Richard A. Harrison, Cristina H. Mandrini,
| Hardi Peter, Saku Tsuneta,
| Bojan Vršnak, Jing-Xiu Wang

TRIENNIAL REPORT 2006-2009

##  1 Introduction

Commission 10 deals with solar activity in all of its forms, ranging from the
smallest nanoflares to the largest coronal mass ejections. This report reviews
scientific progress over the roughly two-year period ending in the middle of
2008. This has been an exciting time in solar physics, highlighted by the
launches of the Hinode and STEREO missions late in 2006. The report is
reasonably comprehensive, though it is far from exhaustive. Limited space
prevents the inclusion of many significant results. The report is divided into
the following sections: Photosphere and Chromosphere; Transition Region;
Corona and Coronal Heating; Coronal Jets; Flares; Coronal Mass Ejection
Initiation; Global Coronal Waves and Shocks; Coronal Dimming; The Link Between
Low Coronal CME Signatures and Magnetic Clouds; Coronal Mass Ejections in the
Heliosphere; and Coronal Mass Ejections and Space Weather. Primary authorship
is indicated at the beginning of each section.

##  2 Photosphere and Chromosphere (C. J. Schrijver)

###  2.1 Quiet-Sun Field Within the Photosphere

The Solar Optical Telescope (SOT) on board the Hinode spacecraft provides an
unprecedented combination of spatial resolution and continuity of
observations. The SpectroPolarimeter focal-plane instrumentation exploits that
to measure the polarization signals from the photospheric plasma. Lites et al.
(2008) and Ishikawa et al. (2008) show direct evidence that much of the
magnetic field in the quiet-Sun photosphere is essentially horizontal to the
solar surface. This observation is the direct confirmation of the existence of
the weak field for which less direct evidence had been found by Harvey et al.
(2007), and contributes to Hanle de-polarization effects discussed by, e.g.,
Trujillo Bueno et al. (2004).

The nearly vertical component is found primarily in the downflow network of
the granular convection, corresponding to the well-known network field. The
horizontal field is mostly found in the interior of the convective cells.
Despite this significant preference for a separation by upflow and downflow
domains, flux has been observed to also emerge already largely vertical even
within the interior of the granular convective cells (Orozco-Suarez et al.
2008). And, perhaps not surprisingly, the conceptually expected evolutionary
pattern of emerging flux is also seen: Centeno et al. (2007) report on
observations in which field is seen to first surface nearly horizontally and
subsequently—as it is advected to the downflow lanes—rights itself to be
nearly vertical.

Lites et al. (2008) measure the mean flux density of the horizontal field to
be about five times higher than that associated with the nearly vertical field
component. Interestingly, radiative MHD simulations of near-surface stratified
convection by Schuessler and Voegler (2008) show a very similar orientation-
dependent ratio for the field. Steiner et al. (2008) reach a similar
conclusion based on their numerical experiments: they argue that the granular
upflows allow field to be stretched horizontally, being advected from over the
cell centers only slowly in the stagnating, overshooting, upper-photospheric
flows. Both studies support the conclusion that near-surface turbulent dynamo
action significantly contributes to the internetwork photospheric field. A
study by Abbett (2007) elucidates how such a turbulent-dynamo field would
connect sub-photospheric and coronal layers through a complex and dynamic
chromospheric layer in between; work by Isobe et al. (2008) explores
numerically the frequent reconnective interactions expected with the overlying
chromospheric canopy field, suggesting that this and the associated wave
generation may have significant consequences for atmospheric heating and
driving of the solar wind.

###  2.2 The Solar Dynamo(s): Global and Local Aspects

Ephemeral bipolar regions are at the small end of the active region spectrum.
Their properties over the solar cycle are an extension of those of their
larger counterparts: they follow the general butterfly pattern, and have the
proper preferential orientation of their dipole axes relative to the equator,
but with a spread about the mean that increases towards the smaller bipoles.
In this, they are a natural extension of the active region population. Where
they were known to differ from the large regions is in the fact that they are
the first to appear and last to fade for a given sunspot cycle.

Now, work by Hagenaar et al. (2008) uncovers another distinct property of
ephemeral regions: the emergence frequency decreases with increasing local
flux imbalance (consistent with findings by Abramenko et al. (2006) and Zhang
et al. (2006) who differentiated only coronal hole regions from other quiet-
Sun regions). Hagenaar et al. (2008) find that the rate of flux emergence is
lower within strongly unipolar network regions by at least a factor of 3
relative to flux-balanced quiet Sun. One consequence of this is that because
coronal holes overlie strong network regions, there are fewer ephemeral
regions, and therefore fewer EUV or X-ray bright points within coronal holes.

The ephemeral-region population thus takes an interesting position in the
study of solar magnetic activity: with the smallest-scale internetwork field
perhaps largely generated by a local turbulent dynamo (Schuessler & Voegler
2008), and with the active regions associated with a global dynamo action, the
ephemeral region population has signatures of both. Voegler & Schuessler
(2007) show that local dynamo action can lead to a mixed-polarity field
similar to the flux balanced very-quiet network field. It remains to be seen
what such experiments predict in case there is a net flux imbalance, i.e., a
background ‘guide field’: are fewer ephemeral regions generated, or does
reconnection with the background guide field cause fewer of them to survive
the rise to the surface (see discussion by Hagenaar et al. 2008).

###  2.3 Emerging Flux: Observations and Numerical Experiments

Observations made with the Hinode SOT show unambiguously that magnetic flux
bundles that form active regions do not emerge as simply curved arches, but
rather as fragmented collections of undulating flux bundles. Each bundle
likely crosses the photosphere one or more times between the extremes of the
emerging region (e.g., Lites 2008). This is likely the result of the coupling
to the near-surface convective motions, and the difficulty of relatively heavy
sub-photospheric material to drain from the dipped field segments.
Reconnection between neighboring supra-photospheric flux bundles could pinch
off the sub-photospheric mass pockets, thus allowing the field to rise into
the corona. Radiative MHD simulations of emerging flux by Cheung et al. (2008)
support this interpretation: they show the ‘serpentine’ nature of the emerging
flux, with characteristics that resemble the observed patterns of emerging
flux, flux cancellation with associated downflows, convective collapse into
strong-field flux concentrations, and photospheric bright points. Note that an
example of field dipping into sub-photospheric layers is also discussed by
Abbett (2007).

###  2.4 Upper-Chromospheric Dynamics: Spicules and Waves

Spectacularly sharp Ca II H narrow-band filter observations made both with
Hinode’s SOT and the Swedish Vacuum Solar Telescope reveal ubiquitous jet-like
features (called spicules or fibrils) above the solar limb. The relatively
long-lived, broad population among these (discussed by De Pontieu et al.
2007a) appear to be caused by acoustic shock waves propagating upward from the
photosphere. These shock waves cause the chromospheric material to undulate
with almost perfectly parabolic height-time profiles, and saw-tooth velocity
patterns. These shock-induced fibrils occur both in plages and in quiet Sun.

A more enigmatic phenomenon is formed by the much finer and more transient
population of hair-like high extension of the chromospheric plasma discussed
by De Pontieu et al. (2007b). Their origin remains subject to debate, but
their transverse displacements point to the ubiquitous existence of Alfven-
like waves propagating into the corona. This is the most direct observational
evidence for the existence of such waves reported to date. The estimated power
suffices to heat the quiet-Sun corona and power the solar wind: these waves
have amplitudes of 10-25 km s-1 for periods of 100-500 seconds. Alfven-like
waves with similar periods have also been observed for the first time in
coronal loops, using the CoMP instrument (Tomczyk et al. 2007).

##  3 Transition Region (H. Peter)

The transition region from the chromosphere to the corona, originally thought
of as a simple thin onion shell-like layer, is a spatially and temporally
highly complex part of the solar atmosphere. So far we are missing a unifying
picture combining the numerous phenomena observed in emission lines formed
from a couple of 10,000 K to several 100,000 K. Some of the aspects are re-
interpreted by Judge (2008), who attempts to explain the transition region as
being due to cross-field conduction of neutral atoms. The emission measure
increasing towards low temperatures and the persistent redshifts are two of
the major observational facts to be explained. A collection of one-dimensional
transient models (Spadaro et al. 2006) and a three-dimensional MHD model
(Peter et al. 2006) gave quantitative explanations for this. In agreement with
the latter model, Doschek (2006) could show that the bulk of the (low)
transition region emission originates from small cool loops. Using rocket
imaging data, Patsourakos et al. (2007) give further direct observational
evidence for the existence of such small loop-like structures dominating the
transition region emission.

The magnetic field has to connect the transition region structures down to the
chromosphere, and in case they are part of hot coronal elements also to the
corona (Peter 2007). However, correlations between the transition region and
the photosphere cannot be identified in a unique way (Sánchez Almeida et al.
2007). Larger features in the photosphere, such as moving magnetic features,
might well leave an imprint in the outer atmosphere (Lin et al. 2006). In
general, the connection from the chromosphere to the transition region is
quite subtle and hard to identify in observations (Hansteen et al. 2007). A
new way to investigate the relation between transient events in the transition
region and the chromosphere was presented by Innes (2008). She studied the
chromospheric emission of molecular hydrogen near 111.9 nm during microflaring
events and proposes that the (coronal) energy deposition in the microflare
also heats the chromosphere and thus affects the opacity for molecular
hydrogen lines.

The energy balance, one part being the heating process, is largely determining
the pressure of the transition region and thus implicitly also the mass
loading. Combining various models and observations, Aschwanden et al. (2007)
argue that the bulk part of the heating is located deep down, basically
reflecting an exponential decay of the heating rate with height, on average.
As speculated earlier, Tian et al. (2008a) could now show that the persistent
blueshifts in upper transition region lines are not due to the solar wind
outflow, but due to mass loading of loops. Recent Doppler shift observations
with the Extreme ultraviolet Imaging Spectrometer (EIS; Culhane et al. 2007a)
on board Hinode indicate that the redshifts are due to radiative cooling and
subsequent bulk downflows within the loops (Bradshaw 2008).

New investigations of coronal moss, i.e. the (upper) transition region
footpoint areas of large hot loops, show that in moss regions the temperature
is inversely related to the density (Tripathi et al. 2006). Using comparisons
with models, Warren et al. (2008b) show that in order to understand this moss
within the framework of a steady uniform heating model, one needs to assume
that the moss plasma is not fully filling the volume. However, it remains to
be seen if such a static model is applicable at all, because one might suspect
a spatially varying heating rate (see above), and if the assumption of static
moss is justified.

Motivated by direct magnetic field measurements in the corona indicating the
presence of Alfven waves (Tomczyk et al. 2007) and observations with SOT on
Hinode above the limb, McIntosh et al. (2008) re-interpreted the widths of
transition region lines across the solar disk. They conclude that the present
observations are consistent with a line-of-sight superposition of Alfvenic
disturbances in small-scale structures. How this relates to the new finding of
Doschek et al. (2007) that the (non-thermal) line widths are largest not in
the brightest parts of an active region but in dimmer regions adjacent to
bright loops remains to be seen. Doschek et al. (2007) find broad lines
related to potential outflow locations, so maybe this problem hints to
different acceleration and heating mechanisms in open and closed field
regions. Another new difference between (globally) open and closed field
regions, was proposed by Tian et al. (2008b), who find evidence that the
expansion of transition region structures is more rapid in the coronal holes
as compared to the quiet Sun. Dolla & Solomon (2008) analyzed line widths
above the limb in order to determine the (kinetic) temperatures of minor ions
in presumably open field regions. They find the smallest mass-to-charge-ratio
ions to be the hottest at a given height, but their analysis remains
inconclusive with regard to supporting or disproving the proposed heating by
ion-cyclotron resonances.

While being in orbit nearly 13 years now, the instruments on board SOHO, SUMER
in particular, still give numerous new valuable results on the transition
region. The EIS instrument on board Hinode covers wavelengths around 17-21 nm
and 24-29 nm. This mainly includes emission lines formed from 1 to several MK,
but also a small number of lines from the transition region, and allows good
density diagnostics (Feldman et al. 2008). Given the spectral range, the main
science topics are grouped around active region phenomena, while the
transition region can also be investigated (Young et al. 2007). Besides these
instruments, which will provide the main source for observations of transition
region lines in the coming years, rocket experiments complement these data.

##  4 Corona and Coronal Heating (J. A. Klimchuk)

The past two years have seen considerable progress in understanding the
magnetically-closed corona and how it is heated. This short report highlights
just some of the important contributions. Much effort has been devoted to
determining the properties of the heating—how it varies in time and space and
whether it depends on physical parameters such as the strength of the magnetic
field and the length of field lines. Some studies have concentrated on
individual coronal loops, while others have addressed active regions as a
whole. These efforts have both clarified some issues and raised new questions.

Let us first consider distinct, measurable loops. A short history is useful.
For many years after the Skylab soft X-ray observations, it was thought that
loops are static equilibrium structures maintained by steady heating. Then
came the EUV observations from EIT/SOHO and TRACE. These revealed that warm
($\sim$ 1 MK) loops are much too dense for static equilibrium and have super-
hydrostatic scale heights. Modeling efforts showed that the excess densities
and large scale heights could be explained by impulsive heating. Because of
their temperature response, EIT and TRACE are sensitive to the loops when they
are cooling by radiation, well after the heating has ceased. The problem is
that loops are observed to persist for longer than a cooling time, so a
monolithic model is not viable. This led to the suggestion that loops are
bundles of thin, unresolved strands. The observed high densities, large scale
heights, and long lifetimes can all be explained if the strands are heated at
different times by a storm of nanoflares. Since the strands are in different
stages of cooling, a range of temperatures should be present within the loop
bundle at any given time. In particular, there should be a small amount of
very hot ($>5$ MK) plasma. See Klimchuk (2006) for a discussion of these
points and original references.

Whether loops are isothermal or multi-thermal has been intensely debated over
the past several years. Double- and triple-filter observations from TRACE seem
to suggest that the most narrow loops are isothermal (Aschwanden 2008).
However, it has been demonstrated that many different thermal distributions,
including ones that are broad, can reproduce the observed intensities, even
with 3 filters (Schmelz et al. 2007a; Patsourakos & Klimchuk 2007; Noglik &
Walsh 2007). Spectrometer observations provide far superior plasma
diagnostics. The results here are mixed. Studies made with CDS/SOHO continue
to find evidence for both isothermal loops and highly multi-thermal loops
(Schmelz et al. 2007b; Cirtain et al. 2007), while studies made with the new
EIS instrument on Hinode find that loops tend to be mildly multi-thermal
(Ugarte-Urra et al. 2008; Warren et al. 2008a). Where temporal information is
available, there is clear evidence that the loops are evolving, but the
evolution is generally slower than expected for radiative cooling. Loop
lifetimes are extremely important and require further investigation. A loop
bundle will be only mildly multi-thermal if the storm of nanoflares is short-
lived; however, the observed lifetime of the loop will then be correspondingly
short. If a loop is observed to persist for much longer than a cooling time
(López Fuentes et al. 2007), then its thermal distribution is expected to be
broad. More work is needed on whether the lifetimes and thermal distributions
of loops are consistent. Finally, Landi & Feldman (2008) have found that one
particular active region is dominated by three distinct temperatures, which
would greatly challenge our understanding if correct.

Modeling the plasma properties of whole active regions is a relatively new
endeavor. In addition to providing valuable information on coronal heating,
these research models are the forerunners of eventual operational models for
nowcasting and forecasting the solar spectral irradiance. This is of great
practical value, since the irradiance controls the dynamics, chemistry, and
ionization state of the terrestrial upper atmosphere and thereby affects radio
signal propagation, satellite drag, etc. Active region models based on static
equilibrium are able to reproduce the observed soft X-ray emission reasonably
well, but they fail, often miserably, at reproducing the EUV emission. The
model corona is too faint in the EUV (there are no warm loops) and the moss at
the transition region footpoints of hot loops is too bright. Winebarger et al.
(2008) have demonstrated that better agreement can be obtained in the core of
an active region by using a combination of flux tube expansion and filling
factors near 10%. Filling factors of this magnitude have been measured in moss
with EIS (Warren et al. 2008b). Small filling factors are consistent with the
idea of unresolved loop strands. Reale et al. (2007) have developed a new
multi-filter technique using XRT/Hinode data that reveals considerable thermal
structure on small but resolvable scales.

Active region models based on impulsive heating are in much better agreement
with observations than are static models. In particular, the predicted coronal
EUV emission is greatly enhanced (Warren & Winebarger 2007; Patsourakos &
Klimchuk 2008a). The predicted moss emission is still too bright, but these
models assume constant cross section flux tubes, and expanding tubes will
improve the agreement, as they do for static models. It is also likely that
the brightness of the observed moss is diminished by spicules and possibly by
other cool absorbing material.

Nanoflare models predict that very hot plasma should be present throughout the
corona, albeit in very small quantities (Klimchuk et al. 2008; this paper
presents a highly efficient IDL code for modeling dynamic loops and is
available upon request). The intensities of very hot spectral lines are
expected to be extremely faint, due to the small emission measures and
possible also to ionization nonequilibrium effects (Bradshaw & Cargill 2006;
Reale & Orlando 2008). Measureable quantities of very hot ($\sim$ 10 MK)
plasma have been detected outside of flares by the CORONAS-F
spectroheliometers (Zhitnik et al. 2006), RHESSI (McTiernan 2008), and XRT
(Siarkowski et al. 2008; Reale et al. 2008). The derived differential emission
measure distributions, DEM($T$), are consistent with the predictions of
nanoflare models. The DEM($T$) derived from EIS spectra for $T\leq 5$ MK are
also consistent with the predictions (Patsourakos & Klimchuk 2008b). Other
tests of the nanoflare idea include emission line Doppler shifts, broadening,
and wing enhancements that are associated with evaporating and condensing
plasma (Patsourakos & Klimchuk 2006; Hara et al. 2008; Bradshaw 2008).

Information on the distribution of nanoflare energies can be inferred from the
intensity fluctuations of observed loops (Parenti et al. 2006; Pauluhn &
Solanki 2007; Parenti & Young 2008; Sarkar & Walsh 2008; Sakamoto et al. 2008;
Bazarghan et al. 2008). Proper flares are known to have a power law energy
distribution with an index $<2$. Extrapolating to smaller energies implies
that nanoflares cannot heat the corona, as first pointed out by Hudson (1991).
However, it is now believed that the power law index for small events is $>2$,
though with a large uncertainty (Benz 2004; Pauluhn & Solanki 2007; Bazarghan
et al. 2008). Furthermore, the subset of proper flares that are not associated
with CMEs also have a power law index $>2$ (Yashiro et al. 2006; see Section
6). Since the physics of eruptive events and noneruptive events (nanoflares
and confined flares) is likely to be much different, it is not surprising that
they obey different power laws.

Thermal nonequilibrium, a phenomenon thought to be important for prominence
formation (Karpen & Antiochos 2008), may also play an important role in
ordinary loops. Loop equilibria do not exist for steady heating if the heating
is highly concentrated low in the loop legs. Instead, cool condensations form
and fall to the surface in a cyclical pattern that repeats on a time scale of
hours. Resolvable condensations are indeed observed in active regions, but
only in a small fraction of loops (Schrijver 2001). As a possible explanation
for other EUV loops, Klimchuk & Karpen (2008) appeal to the multi-strand
concept. The individual tiny condensations that occur within each strand will
not be detected as long as the strands are out of phase. It is encouraging
that the models predict excess densities similar to those of observed EUV
loops. Mok et al. (2008) report thermal nonequilibrium behavior in their
active region simulations. Hot loops form and cool, but without producing
localized condensations.

We can summarize the state of understanding as follows. Much of the
magnetically-closed corona is certainly not in static equilibrium, but much of
it could be. A significant portion—perhaps the vast majority—is heated
impulsively or is in thermal nonequilibrium, or some combination thereof. Most
coronal heating mechanisms that have been proposed involve impulsive energy
release (Klimchuk 2006; Uzdensky 2007; Cassak et al. 2008; Dahlburg et al.
2008; Rappazzo et al. 2008; Ugai 2008). It should be noted, however, that
nanoflares that recur sufficiently frequently within the same flux strand (on
a time scale much shorter than the cooling time) will produce quasi-static
conditions. It is clear that more observational and theoretical work is
required before the coronal heating problem will be solved.

##  5 Coronal Jets (L. van Driel-Gesztelyi)

Coronal bright points are often observed to have jets—collimated transient
ejections of hot plasma. Hinode (Kosugi et al. 2007) can now study the fine
detail of jets which tend to occur preferentially inside coronal holes, which
is consistent with reconnection taking place between the open magnetic field
of the coronal hole and the closed loop field lines. Observations with the XRT
instrument (Golub et al. 2007) revealed that jets from polar coronal holes are
more numerous than previously thought (60 jets day-1, Savcheva et al. 2007;
and even 10 jets hour-1, Cirtain et al. 2007). The EIS instrument (Culhane et
al. 2007a) allows direct measurement of the velocity of jets in the corona for
the first time. The footpoints of the loops are seen to be red-shifted which
is consistent with downflowing cooling plasma following reconnection. The
(blue-shifted) jet is the dominant feature in velocity space but not in
intensity (Kamio et al. 2007). Another new feature of jets is post-jet
enhancement of cooler coronal lines observed by EIS. This can be explained by
the hot plasma in the jet not having sufficient velocity to leave the Sun and
then falling back some minutes later (Culhane et al 2007b).

XRT observations of jets at the poles have shown mean velocities for jets of
160 km s-1 (Savcheva et al. 2007). Multiple velocity components were found in
jets by Cirtain et al. (2007) in XRT polar coronal hole data: a spatio-
temporal average of about 200 km s-1 as well as a much higher velocity
measured at the beginning of each jet—with speeds reaching 800 km s-1. Cirtain
et al. (2007) interpret this early (and sometimes recurrent) fast flow as
being due to plasma ejected at the Alfven speed during the relaxation phase
following magnetic reconnection. The mass flux supplied by about 10 jets per
hour occurring in the two polar coronal holes was estimated to produce a net
flux of 1012 protons m-2 s-1 which is only a factor of 10 less than the
current estimates of the average solar wind flux. These small jets are
providing a substantial amount of mass that is being carried into
interplanetary space. A 3D numerical simulation has been carried out to
compare with these observations (Moreno-Insertis et al. 2008) and is found to
be consistent with several key observational aspects of polar jets such as
their speeds and temperatures.

A study of the 3D morphology of jets became possible for the first time with
stereoscopic observations by the EUVI/SECCHI imagers (Howard et al. 2008)
onboard the twin STEREO spacecraft. The most important geometrical feature of
the observed jets was found to be helical structures showing evidence of
untwisting (Patsourakos et al. 2008). This is in agreement with the 3D model
proposed by Pariat et al. (2008) with magnetic twist (untwisting) being the
jet’s driver.

##  6 Flares (L. Fletcher and J. Wang)

In this brief review we focus on progress in flare energy build-up and flare
prediction, flare photospheric effects, high energy coronal sources, non-
thermal particles, the flare-CME relationship, and recent advances with
Hinode.

Where is the magnetic free energy stored in a flaring active region? Using the
increasingly robust methods for extrapolating magnetic fields from vector
magnetic field measurements, Schrijver et al. (2008) find evidence for pre-
flare filamentary coronal currents located $<20$ Mm above the photosphere and
Regnier & Priest (2007) show that in a newly-emerged active region the free
energy is concentrated within the first 50 Mm (in an older, decaying region it
resides at higher altitudes). Horizontal shear flows close to the neutral line
prior to large flares (Deng et al. 2006) confirm the concentration of free
energy in a small spatial scale, and Schrijver (2007) finds that if the
unsigned flux within 15 Mm of the polarity inversion line exceeds
$2\\!\times\\!10^{21}$ Mx a major flare will occur within a day. Though Leka &
Barnes (2007) find that the probability of flaring has only a weak
relationship to the state of the photospheric magnetic field at any time,
single or synthesized magnetic parameters are being used with some success to
quantify flare probability and productivity. Georgoulis & Rust (2007)
introduce the effective connected magnetic field of an active region, finding
that this exceeds 1600 and 2100 G for M- and X-class flares, respectively, at
95% probability. LaBonte et al. (2007) surveyed the helicity injection prior
to X-class flares producing a CME, finding occurrence only if the peak
helicity flux exceeds $6\\!\times\\!10^{36}$ Mx s-1. Cui et al. (2007) find
that flare probability increases with active region complexity,
nonpotentiality, and length of polarity inversion line. Impulsive phase HXR
sources are concentrated where the magnetic field is strong, and where the
reconnection rate is high (Temmer et al. 2007; Jing et al. 2008; Liu et al.
2008).

The last three years have seen effort directed towards understanding the
magnetic and seismic effects of flares near the photosphere. It is clear that
the photospheric magnetic field changes abruptly and non-reversibly during the
flare impulsive phase (see e.g. Sudol & Harvey 2005 for a recent survey).
Rapid changes in sunspot structure have also been detected by Chen et al.
(2007) in 40% of X-class flares, 17% of M flares, and 10% C flares, while Wang
(2006) finds variations in magnetic gradient close to the polarity inversion
line consistent with a sudden release of magnetic shear. The obvious future
task is to analyze vector magnetograms to identify changes in the ‘twist’
component of the field.

Flare-generated seismic waves, discovered by Kosovichev & Zharkova (1998) and
amply confirmed in Cycle 23, also show the flare’s photospheric impact. Donea
et al. (2006), Kosovichev (2006, 2007) and Zharkova & Zharkov (2007) show that
flare HXR sources and seismic sources correlate in space and time. Seismic
sources are associated also with white light kernels, responsible for the
majority of the flare radiated energy and strongly correlated with HXR sources
(Fletcher et al. 2007) but the total acoustic energy is a small fraction of
total flare energy (Donea et al. 2006). Nonetheless, looking at the cyclic
variation of the total energy in the Sun’s acoustic spectrum, Karoff &
Kjeldsen (2008) propose that—analogous with earthquakes—flares may excite
long-duration global oscillations.

The RHESSI mission has discovered several new types of flare coronal HXR
sources, and we highlight here a hard-spectrum HXR source at least 150 Mm
above the photosphere, with a nonthermal electron fraction of about 10%
(Krucker et al. 2007b). Hard spectrum gamma-ray (200-800 keV) coronal sources
have also been found, suggesting coronal electron trapping (Krucker et al.
2008). A soft-hard-soft spectral variation with time is present in some
coronal sources as well as footpoints (Battaglia & Benz 2006), and this may be
explicable by a combination of coronal trapping and stochastic electron
acceleration (Grigis & Benz 2006). A curious observation with the Owens Valley
Solar Telescope of terahertz emission may come from a compact coronal source
of electrons at 800 keV, if the electrons are radiating in a volume with a
magnetic flux density of 4.5 kG (Silva et al. 2007).

Discovering the origin and properties of the flare electron distribution
continues to motivate advanced modeling and observations. Particle-in-cell
simulations, used for some years in magnetospheric physics, have been
harnessed to study acceleration in coronal magnetic islands produced by
magnetic reconnection (Drake et al. 2006), and wave-particle distributions in
current sheet and uniform magnetic field geometries (Karlicky & Barta 2007;
Sakai et al. 2006). Flare Vlasov simulations are also being developed (e.g.
Miteva et al. 2007; Lee et al. 2008). Detailed RHESSI HXR spectroscopy has led
to a new diagnostic for the flare electron angular distribution, based on
photospheric HXR albedo (Kontar et al. 2006a). This diagnostic suggests that
electron distributions might not be strongly downward-beamed in the
chromosphere (Kontar et al. 2006b; Kasparova et al. 2007; though see Zharkova
& Gordovskyy 2006 for an alternative explanation). Xu et al. (2008) studied
RHESSI flares having an extended coronal source, finding evidence for an
extended coronal accelerator. Looking at the larger coronal context, Temmer et
al. (2008) show that peaks in the flare electron acceleration rate and in the
CME acceleration rate are simultaneous within observational constraints
($\sim$ 5 minutes). RHESSI and WIND observations suggest that the spectral
indices of flare and interplanetary electrons are correlated, but in a way
that is inconsistent with existing models for flare X-ray emission, and the
number of escaping electrons is only about 1/500th of the number of electrons
required to produce the chromospheric HXR flux (Krucker et al. 2007a).

However, outstanding questions about the total electron number and supply in
solar flares has prompted various authors to suggest alternatives to the
‘monolithic’ coronal electron beam picture. For example Dauphin (2007),
Gontikakis et al. (2007) examine acceleration in multiple, distributed coronal
region sites, while Fletcher & Hudson (2008) investigate the transport of
flare energy to the chromosphere in the form of the Poynting flux of large-
scale Alfvenic pulses, in strong and low-lying coronal magnetic fields.

The relationship between flares and CMEs continues to be an important topic.
The general consensus regarding the spatial correspondence between CME
position angle and flare location in the pre-SOHO era was that the flare is
located anywhere under the span of the CME (e.g., Harrison 2006). However,
using 496 flare-CME pairs in the SOHO era, Yashiro et al. (2008) found that
the offset between the flare position and the central position angle (CPA) of
the associated CME has a Gaussian distribution centered on zero, meaning the
flare is typically located radially below the CME leading edge. This finding
suggests a closer flare-CME relationship as implied by the CSHKP eruption
model. Many flares are not associated with CMEs. Yashiro et al. (2006) studied
two sets of flares one with and the other without CMEs. The number of flares
as a function of peak X-ray flux, fluence, and duration in both sets followed
a power law. Interestingly, the power law index was $>2$ for flares without
CMEs, while $<2$ for flares with CMEs. In flares without CMEs, the released
energy seems to go entirely into heating, which suggests that nanoflares may
contribute significantly to coronal heating (see Section 4).

The launch of Hinode promises significant advances in flare physics in the
next cycle. Thus far there have only been a small number of well-observed
large flares, but observations of small flares start to show the combined
power of RHESSI and Hinode. For example, Hannah et al. (2008) find a
microflare not conforming to the usual relationship between flare thermal and
nonthermal emission, and Milligan et al. (2008) show evidence for hot
downflowing plasma in the flare corona, not explained in any existing flare
model. Observational evidence for a new kind of reconnection, called slip-
running reconnection, has been found by Aulanier et al. (2007), and sub-
arcsecond structure in the white light flare sources has been demonstrated by
Isobe et al. (2007). We look forward to the continued operation of these
instruments, and the theoretical advances that they will bring, in the rise of
Cycle 24.

##  7 Coronal Mass Ejection Initiation (L. van Driel-Gesztelyi)

In recent years our physical understanding of CMEs has evolved from cartoons
inspired by observations to full-scale numerical 3D MHD simulations
constrained by observed magnetic fields. Notably, there has been progress made
in simulating CME initiation by flux rope instabilities as inspired by
observed filament motions during eruption which frequently include helical
twisting and writhing (e.g. Rust & LaBonte 2005; Green et al. 2007). Several
of these simulations use the analytical model of a solar active region by
Titov & Démoulin (1999) as initial condition. The model contains a current-
carrying twisted flux rope that is held in equilibrium by an overlying
magnetic arcade. The two instabilities considered as eruption drivers are the
ideal MHD helical kink and torus instabilities. The helical kink instability
sets in if a certain threshold of (flux rope) twist ($\sim$ 2.5 turns for
line-tied flux ropes) is reached (e.g., Török & Kliem 2005). Above this
threshold, twist becomes converted into writhe during the eruption, deforming
the flux rope (or filament) into a helical kink shape. On the other hand, a
current-carrying ring (or flux rope) situated in an external poloidal magnetic
field ($B_{ex}$) is unstable against radial expansion when the Lorentz self-
force or hoop force decreases more slowly with increasing ring radius than the
stabilizing Lorentz force due to $B_{ex}$. Known as the torus instability, its
possible role in solar eruptions has been examined by Kliem & Török (2006) and
Isenberg & Forbes (2007). In solar eruptions, the torus instability does not
require a pre-eruptive, highly twisted flux rope, but (i) a sufficiently steep
poloidal field decrease with height above the photosphere and (ii) an
(approximately) semi-circular flux rope shape. Both the helical kink and torus
instabilities may be responsible for initiating and driving
prominence/filament eruptions and thus CMEs. The magnetic field decrease with
height above the filament was shown to be critical whether a confined eruption
or a full eruption occurs as well as for determining the acceleration profile,
corresponding to fast CMEs for rapid (field) decrease, as it is typical of
active regions, and to slow CMEs for gentle decrease, as is typical of the
quiet Sun (Török & Kliem 2005, 2007; Liu 2008). The latter means that CMEs
from complex active regions with steep field gradients in the corona are more
likely to give rise to fast CMEs—something that is indeed observed.

More complex CME initiation models involve multiple magnetic flux systems,
such as in the magnetic break-out model (e.g. Antiochos et al. 1999, DeVore &
Antiochos 2008). In this model, magnetic reconnection removes unstressed
magnetic flux that overlies the highly stressed core field and this way allows
the core field to erupt. The magnetic break-out model involves specific
nullpoints and separatrices. A multi-polar configuration was also included in
the updated catastrophe model (Lin & van Ballegooijen 2005), the flux
cancellation model (Amari et al. 2007), and the MHD instability models (Török
& Kliem 2007). In an attempt to test CME initiation models with special
attention to the breakout, Ugarte-Urra et al. (2007) analyzed the magnetic
topology of the source regions of 26 CME events using potential field
extrapolations and TRACE EUV observations. They found only seven events which
could be interpreted in terms of the breakout model, while a larger number of
events (12) could not be interpreted in those terms. The interpretation of the
rest remained uncertain. On the other hand, the CME event analyzed by Williams
et al. (2005) provided a good example to indicate that also a combination of
several mechanisms, e.g. magnetic break-out and kink instability, can be at
work in initiating CMEs.

##  8 Global Coronal Waves and Shocks (B. Vršnak)

The research on globally propagating coronal disturbances (large-amplitude
waves, shocks, and wave-like disturbances) continued to be very dynamic. Maybe
the most prominent characteristic of the past triennium was an enhanced effort
to combine detailed multi-wavelength observations with the theoretical
background. The empirical research resulted in a number of new findings,
leading to new ideas and interpretations, whereas the theoretical research
provided a better understanding of physical processes governing the formation
and propagation of global coronal disturbances.

For the first time the EUV signatures of a global coronal wave (‘EIT waves’)
were measured at high cadence by EUVI/STEREO, related to the eruption of 2007
May 19 (Long et al. 2008; Veronig et al. 2008). Long et al. (2008) reported
for the first time the wave signatures at 304 Å. Furthermore, they confirmed
the idea by Warmuth et al. (2001) that velocities of EIT waves measured by
EIT/SOHO are probably significantly underestimated due to the low cadence of
the EIT instrument. Veronig et al. (2008) revealed reflection of the wavefront
from the coronal hole boundary, indicating that the observed disturbance
represents a freely-propagating MHD wave.

The data from pre-STEREO instruments continued to be exploited fruitfully.
Mancuso & Avetta (2008) analyzed the UV-spectrum (UVCS/SOHO) of the 2002 July
23 coronal shock, and concluded that the plasma-to-magnetic presure ratio
$\beta$ could be an important parameter in determining the effect of ion
heating at collisionless shocks. Employing the extensive data on CMEs, solar
energetic particle (SEP) events, and type II radio bursts during the SOHO era,
Gopalswamy et al. (2008b) demonstrated that essentially all type II bursts in
the decameter-hectometric (DH) wavelength range are associated with SEP events
once the source location is taken into account. Shen et al. (2007) proposed a
method to determine the shock Mach number by employing the CME kinematics,
type II burst dynamic spectrum, and the extrapolated magnetic field. Analyzing
one Moreton wave that spanned over almost $360^{\rm o}$, Muhr et al. (2008)
revealed two separate radiant points at opposite ends of the two-ribbon flare,
indicating that the wave was driven by the CME expanding flanks. Veronig et
al. (2006) found out that the Moreton/EIT wave segments where the front
orientation is normal to the coronal hole boundary can intrude into the
coronal hole up to 60-100 Mm.

Regarding the nature of global coronal disturbances, some new ideas appeared.
For example, Attrill et al. (2007) attributed EIT waves to successive
reconnections of CME flanks with coronal loops, which could explain the
association of EIT ‘waves’ and shallow coronal dimmings which are formed
behind the bright front. Wills-Davey et al. (2007) proposed that slow EIT
waves are caused by MHD slow-mode soliton-like waves. Delannee et al. (2008)
performed a 3D MHD simulation to show that EIT ‘waves’ could be a signature of
a current shell formed around the erupting structure. Balasubramaniam et al.
(2007) demonstrated that the visibility of Moreton waves increases when
sweeping over filaments and filament channels, so they put forward the idea
that a significant contribution to the Moreton-wave Hα signature might be
coming from coronal material of enhanced density.

The question of the origin of coronal shocks and large amplitude waves
continues to be one of the central topics in this field. The published studies
showed a variety of results, some biased towards the CME-driven option, some
favoring the flare-ignited scenario, and some finding arguments for small
scale ejecta (e.g., Chen 2006; Pohjolainen & Lehtinen 2006; Shanmugaraju et
al. 2006a,b; Subramanian & Ebenezer 2006; Cho et al., 2007; Liu et al. 2007;
Reiner et al. 2007; White 2007; Grechnev et al. 2008; Muhr et al. 2008;
Veronig et al. 2008; Magdalenic et al. 2008; Mancuso & Avetta 2008;
Pohjolainen 2008). To illustrate the current level of ambiguity in such
studies, let us mention that for one well observed event two sets of authors
came to diametrally opposite conclusions: Vršnak et al. (2006) favored a flare
driver, whereas Dauphin et al. (2006) advocate a CME. The status of the
‘CME/flare controversy’ was reviewed recently by Vršnak & Cliver (2008).

Related to the formation and propagation of large-amplitude waves and shocks,
a number of important theoretical papers were published. Pagano et al. (2007)
investigated the role of magnetic fields and showed that a CME-driven wave
propagates to longer distances in the absence of magnetic field than in the
presence a weak open field. Ofman (2007) modeled the wave activity following a
flare by launching a velocity pulse into a model active region and
demonstrated that the resulting global oscillations are in good agreement with
observations. Employing the photospheric magnetic field measurements, Liu et
al. (2008) performed a 3D MHD simulation of a CME, and showed that the shock
segment at the nose of the CME remains quasi-parallel most of the time. In the
simulation of reconnection in a vertical current sheet, Barta et al. (2007)
revealed the formation of large-amplitude waves associated with changes of the
reconnection rate, which might explain flare-associated type II bursts in the
wake of CMEs. Zic et al. (2008) developed an analytical MHD model describing
the formation of large-amplitude waves by impulsively expanding 3D pistons.
The model provides an estimate of the time/distance at which the shock should
be formed, dependent on the source-surface acceleration, the terminal
velocity, the initial source size, the ambient Alfven speed, and plasma
$\beta$.

Finally, it should be noted that a comprehensive review on coronal waves and
shocks was published by Warmuth (2007). Gopalswamy (2006e) reviewed the
relationship between CMEs and type II bursts, while Mann & Vršnak (2007)
surveyed the relationship between CMEs, flares, coronal shocks, and particle
acceleration.

##  9 Coronal Dimming (R. Harrison and L. van Driel-Gesztelyi)

There is no strict definition of the phenomenon which we call coronal dimming.
Most authors consider coronal dimming to be a depletion of extreme-UV (EUV) or
X-ray emission from a large region of the corona, which is thought to be
closely associated with coronal mass ejection (CME) activity. Clearly,
understanding the onset phase of a CME is one of the key issues in solar
physics today, so the study of such dimming activity could well be of critical
importance. However, most of the literature deals with dimming in a rather
hand-waving manner, with the emphasis on phenomenological studies and
associations, no strict definitions of what constitutes a dimming event (e.g.,
the depth of the depletion in intensity, the size of the dimming region, etc.)
and little in terms of a physical interpretation of the plasma characteristics
of the dimming region. Having said that, some key studies are emerging which
do tackle such issues head on, and with the advent of the new STEREO and
Hinode spacecraft, along with the on-going SOHO and TRACE missions, as well as
the up-coming SDO mission, we have many tools to address this area of research
effectively.

Coronal dimming is not a newly discovered phenomenon; Rust and Hildner (1976)
reported such an event using Skylab observations. More recently, from the late
1990s, dimming was reported using X-ray and EUV, imaging and spectroscopic
data, from the SOHO and Yohkoh spacecraft (e.g. Sterling & Hudson 1997;
Harrison 1997; Gopalswamy & Hanaoka 1998; Zarro et al. 1999; Harrison & Lyons
2000), and dimming has taken center-stage in the study of mass ejection onset
in recent years (e.g., recent studies include Moore & Sterling 2007; Zhang et
al. 2007; Reinard & Biesecker 2008). In many ways coronal dimming has become a
well established phenomenon.

The majority of dimming reports involve EUV or X-ray imaging, and we have
excellent tools aboard SOHO, TRACE, STEREO and Hinode to identify and study
the topology and evolution of dimming regions. On the other hand, there are
spectroscopic studies of dimming which are providing key plasma information,
despite having limited fields of view or cadence. The combination of imaging
and spectroscopy is essential, but it is worth stressing some of the
spectroscopic studies because they stress the physical processes which are
involved in the dimming and, perhaps, the CME onset process.

EUV spectroscopy has been used to confirm that the dimming process represents
a loss of mass—i.e., it is a density depletion—rather than a change in
temperature (Harrison & Lyons 2000; Harrison et al. 2003). Indeed, these
studies have demonstrated the loss of between $4.3\\!\times\\!10^{10}$ and
$2.7\\!\times\\!10^{14}$ kg, in each case consistent with the mass of an
overlying, associated CME. If we are identifying the plasma which becomes
(part of) the CME, then this is an exciting phenomenon; studies focusing on
the properties of the dimming plasma, before, during and after the event, will
be essential for understanding the CME onset (Harrison & Bewsher 2007).

Hudson et al. (1996) showed that the timescale of the dimming formation
observed in Yohkoh/Soft X-ray Telescope (SXT; Tsuneta et al. 1991) data is
much faster than corresponding conductive and radiative cooling times. More
recently, data obtained by the Hinode/Extreme ultra-violet Imaging
Spectrometer (EIS; Culhane et al. 2007) have shown detection of Doppler
blueshifted plasma outflows of velocity $\approx 40$ km s-1 corresponding to a
coronal dimming (Harra et al. 2007). This result confirms a similar finding
(Harra & Sterling 2001) obtained with the SOHO/Coronal Diagnostic Spectrometer
(CDS; Harrison et al. 1995). In addition, SOHO/CDS limb observations have been
used to show the formation of a dimming region through the outward expansion
of pre-CME EUV loops (Harrison & Bewsher, 2007), which is consistent with such
blueshifts. Imada et al. (2007) find that Hinode/EIS data of a dimming shows a
dependence of the outflow velocity on temperature, with hotter lines showing a
stronger plasma outflow (up to almost 150 km s-1. These works collectively
support the primary interpretation of coronal dimmings as being due to plasma
evacuation.

Statistical studies are becoming important in truly establishing the
relationship with CMEs. Reinard & Biesecker (2008) have recently studied the
properties of 96 dimming events, using EUV imaging, associated with CME
activity. They confirmed earlier studies which showed that the dimming events
could be long-lasting, ranging from 1 to 19 hours, and compared the size of
the dimming regions to the associated CMEs. They also tracked the number of
dimming pixels through each event and showed that the ‘recovery’ after the
dimming often took the form of a two-part slope (plotted as dimming area vs.
time).

Bewsher et al. (2008) have produced the first statistical and probability
study of the dimming phenomenon using spectroscopy. They recognized that while
we have associated CMEs and dimming, there has not been a thorough statistical
study which can really identify the degree of that association, i.e., to put
that relationship on a firm footing. Using spectroscopy, they also recognized
the importance of studying this effect for different temperatures. They made
use of over 200 runs of a specific campaign using the SOHO spacecraft with an
automated procedure for identifying dimming.

Key results included the following: Up to 84% of the CMEs in the data period
can be back-projected to dimming events—and this appears to confirm the
association that we have been proposing. However, they also showed, as did
other spectral studies, that the degree of dimming varies between temperatures
from event to event. If different dimming events have different effects at
different temperatures then this is a problem for monitoring such events with
fixed-wavelength imagers.

Assuming that magnetic field lines of the CME are mostly rooted in the
dimmings, several properties derived from the study of dimmings can be used to
obtain information about the associated CME. Firstly, calculations of the
emission measure and estimates of the volume of dimmings can give a proxy for
the amount of plasma making up the CME mass (Sterling & Hudson 1997; Harrison
& Lyons 2000; Harrison et al. 2003; Zhukov & Auchère 2004). Secondly, the
spatial extent of coronal dimmings can give information regarding the angular
extent of the associated CME (Thompson et al. 2000; Harrison et al. 2003;
Attrill et al. 2007, van Driel-Gesztelyi et al. 2008). Thirdly, quantitative
measurement of the magnetic flux through dimmings can be compared to the
magnetic flux of modeled magnetic clouds (MC) at 1 AU (Webb et al. 2000;
Mandrini et al. 2005; Attrill et al. 2006; Qiu et al. 2007), see Démoulin
(2008) for a review. Fourth, studying the evolution of the dimmings,
particularly during their recovery phase can give information about the
evolution of the CME post-eruption (Attrill et al. 2006; Crooker & Webb 2006)
providing proof for e.g. magnetic interaction between the expanding CME and
open field lines of a neigboring coronal hole. Finally, study of the
distribution of the dimmings, their order of formation and measurement of
their magnetic flux contribution to the associated CME enabled Mandrini et al.
(2007) to derive an understanding of the CME interaction with its surroundings
in the low corona for the case of the complex 28 October 2003 event. They,
building on the model proposed by Attrill et al. (2007), demonstrated that
magnetic reconnection between field lines of the expanding CME with
surrounding magnetic structures ranging from small- to large-scale (magnetic
carpet, filament channel, active region) make some of the field lines of the
CME ‘step out’ from the flaring source region. Magnetic reconnection is driven
by the expansion of the CME core resulting from an over-pressure relative to
the pressure in the CME’s surroundings. This implies that the extent of the
lower coronal signatures match the final angular width of the CME. Through
this process, structures over a large-scale magnetic area become CME
constituents (for a review see van Driel-Gesztelyi et al. 2008). From the
wide-spread coronal dimming some additional mass is supplied to the CME.

Observations show that coronal dimmings recover whilst suprathermal uni- or
bi-directional electron heat fluxes are still observed at 1 AU in the related
ICME, indicating magnetic connection to the Sun. The questions why and how
coronal dimmings disappear whilst the magnetic connectivity is maintained was
investigated by Attrill et al. (2008) through the analysis of three CME-
related dimming events. They demonstrated that dimmings observed in SOHO/EIT
data recover not only by shrinking of their outer boundaries but also by
internal brightenings. They show that the model developed in Fisk & Schwadron
(2001) of interchange reconnections between ‘open’ magnetic field and small
coronal loops is applicable to observations of dimming recovery. Attrill et
al. (2008) demonstrate that this process disperses the concentration of ‘open’
magnetic field (forming the dimming) out into the surrounding quiet Sun, thus
recovering the intensity of the dimmings whilst still maintaining the magnetic
connectivity of the ejecta to the Sun.

Although this brief summary cannot report on all studies, it is clear that we
have made progress very recently in putting the dimming phenomenon on a firm
footing—the association is real—and we are making in-roads into studies of the
plasma activities leading to the dimming/CME onset process. With the
continuation of the SOHO mission, as well as TRACE, combined with the new
STEREO and Hinode missions and the up-coming SDO mission, this is a topic
which will receive much attention in the next few years.

##  10 The Link Between Low-Coronal CME Signatures and Magnetic Clouds (C.
Mandrini)

A major step to understanding the variability of the space environment is to
link the sources of coronal mass ejections (CMEs) to their interplanetary
counterparts, mainly magnetic clouds (MCs), a subset of interplanetary CMEs
characterized by enhanced magnetic field strength when compared to ambient
values, a coherent and large rotation of the magnetic field vector, and low
proton temperature (Burlaga 1995). Identifying the solar sources and comparing
qualitatively and quantitatively global characteristics and physical
parameters both in the Sun and the interplanetary medium provide useful tools
to constrain models in both environments.

Under the assumption that dimmings (see Section 9) at the Sun mark the
position of ejected flux rope footpoints (Webb et al. 2000), the magnetic flux
through these regions can be used as a proxy for the magnetic flux involved in
the ejection and, thus, be compared to the magnetic flux in the associated
interplanetary MC. Another proxy for the flux involved in an ejection is the
reconnected magnetic flux swept by flare ribbons, as they separate during the
evolution of two-ribbon flares. Using EUV dimmings as proxies and
reconstructing the MC structure from one spacecraft observations, Mandrini et
al. (2005) and Attrill et al. (2006) found that the magnetic flux in dimming
regions was comparable to the azimuthal MC flux, while the axial MC flux was
several times lower. Qui et al. (2007) analyzed and compared the reconnected
magnetic flux to the total MC flux, finding similar results (see also
Yurchyshyn et al. 2006; Longcope et al. 2007; Möstl et al. 2008, where MC data
from two spacecraft were used). These results led to the conclusion that the
ejected flux rope is formed by successive reconnections in a sheared arcade
during the eruption process, as opposed to the classical view of a previously
existing flux rope being ejected. However, in extreme events that occur in not
isolated magnetic configurations, it was found that the flux in dimmings did
not agree with the MC flux (Mandrini et al. 2007). This mismatch led these
authors to propose a scenario in which dimmings spread out to large distances
from the initial erupting region through a stepping reconnection process (in a
similar process to that proposed by Attrill et al., 2007, for the
interpretation of EIT waves). An overview of earlier works on quantitative
comparisons of solar and interplanetary global magnetohydrodynamic invariants,
such as magnetic flux and helicity, can be found in Démoulin (2008).

Qualitative comparisons are also useful tools to understand the eruption
process. Studying the temporal and spatial evolution of EUV dimmings, together
with soft X-ray coronal observations, in conjunction with interplanetary in
situ data of suprathermal electron fluxes, Attrill et al. (2006) and Crooker &
Webb (2006) derive an eruption scenario in which interchange magnetic
reconnection between the expanding CME loops and the open field lines of a
polar coronal hole led to the opening of one leg of the erupting flux rope.
Harra et al. (2007), combining EUV and Hα solar observations of eruptive
events with in situ magnetic field and suprathermal electron data, were able
to understand the sequence of events that produced two MCs with opposite
magnetic field orientations from the same magnetic field configuration.

The simple comparison of the magnetic field orientation in the erupting
configurations, which can be inferred from magnetograms, the directions of
filaments, coronal arcades or loops, with the axis of the associated MCs, can
give clues about the mechanism at the origin of solar eruptions. Green et al.
(2007) analyzed in detail associations of filament eruptions and corresponding
MCs, and they found that when the filament and MC axis differed by a large
angle, the direction of rotation was related to the magnetic helicity sign of
the erupting configuration (see also Harra et al. 2007). The rotation was
consistent with the conversion of twist into writhe, under the ideal MHD
constraint of helicity conservation, providing support for the assumption of a
flux rope topology where the kink instability sets in during the eruption (see
the review by Gibson et al. 2006).

##  11 Coronal Mass Ejections in the Heliosphere (R. Harrison)

In the 1970s the Helios spacecraft operated from solar orbits with perihelion
0.31 AU. Zodiacal light photometers were used to detect CMEs in the inner
heliosphere (see e.g. Richter et al. 1982; Jackson & Leinert 1985). CME images
were constructed from three photometers which scanned the sky using the
spacecraft rotation. More recently, a major advance was made with the launch,
in 2003, of the Solar Mass Ejection Imager (SMEI) aboard the Coriolis
spacecraft (Eyles et al. 2003). This instrument maps the entire sky with three
cameras each scanning $60^{\rm o}$ slices of the sky as the spacecraft moves
around the Earth, and thus, it has pioneered full-sky mapping aimed
specifically at the detection of CMEs propagating through the inner
heliosphere (see e.g. recent papers by Kahler & Webb 2007 and Jackson et al.
2007).

The combination of wide-angle heliospheric mapping from out of the Sun-Earth
line is now being satisfied by the Heliospheric Imagers (HI) (Harrison et al.
2008) aboard the NASA STEREO spacecraft. The development of these instruments
has come very much from the SMEI heritage and, with the unique opportunities
from the STEREO spacecraft locations, these instruments are able to image
those CME events directed towards the Earth. Indeed, for the first time, the
HI instruments provide a view of the passage of CMEs along virtually the
entire Sun-Earth line and such observations represent a major milestone in
investigations of the influence of solar activity on the Earth and human
systems.

Each HI instrument consists of two wide-angle telescopes mounted within a
baffle system enabling imaging of the heliosphere from the corona out to
Earth-like distances and beyond. The low scattered light levels and
sensitivity allow the detection of stars down to magnitudes of 12-13. This
performance is excellent for the detection of solar ejecta and solar wind
structure through the detection of Thomson scattered photospheric light off
free electrons in regions of density enhancement.

The STEREO spacecraft were launched in October 2006 with full scientific
operation of the HI instruments starting from April 2007. The spacecraft are
in near Earth-like solar orbits, with one ahead and one behind the Earth in
its orbit. They are drifting away at $22.5^{\rm o}$ per year (Earth-Sun-
spacecraft angle). The spacecraft are labelled STEREO A and STEREO B, for
ahead and behind.

The first HI observations of CMEs in the heliosphere, tracked out to Earth-
like distances, were reported by Harrison et al. (2008). The same instruments
are also reaping the benefits of wide-angle imaging of the heliosphere with
observations of comets (Fulle et al. 2007; Vourlidas et al. 2007), even the
imaging of co-rotating interaction regions (Sheeley et al. 2008a,b; Rouillard
et al. 2008a) and impacts of CMEs at other planets (Rouillard et al. 2008b).

With the HI instruments we now have a real opportunity to begin to relate the
coronal events that we call CMEs with their heliospheric counterparts,
commonly referred to as ICMEs - Interplanetary CMEs. Most ICME studies have
been performed utilizing in situ particle and field observations, and it is
clear that heliospheric imaging can provide a thorough test of the
interpretation of such in situ data on the topology and propagation of CMEs in
the heliosphere. Indeed, the uniqueness of this opportunity is well
illustrated by the fact that there are a number of extremely basic
observational tests which can be made with the new facility to underline our
current understanding of how CMEs travel out through the Solar System.

Crooker & Horbury (2006) have recently reviewed the propagation of ICMEs in
the heliosphere, utilizing in situ data. They note that cartoon sketches of
ICMEs commonly show magnetic field lines connected to the Sun at both ends.
Furthermore, reporting on the work of Gosling et al. (1987), Crooker et al.
(2002), and others, they note that it is widely accepted that counter-
streaming particle beams in ICMEs are a sign that both ends of the ICME are
indeed connected to the Sun. This is known as a ‘closed’ ICME. On the other
hand, uni-directional beams may signal connection at only one end—an ‘open’
ICME. Logically, then, the lack of beams would appear to signal disconnection
at both ends. In this case the ICME has become an isolated plasmoid.

Given this interpretation, in situ observations of ICMEs appear to show many
events which are apparently connected to the Sun at both footpoints, and
rather fewer events which appear to be connected at one end. Complete
disconnection of an ICME (a plasmoid) appears to be rare. In addition, the in
situ observations suggest that CMEs are connected to the Sun over extremely
long distances; Riley et al. (2004) looked for the degree of ‘openness’ of
ICMEs using observations of counter-streaming electrons from Ulysses data and
could detect no trend in the openness of ICMEs with distance out to Jupiter.
If ICME connectivity to the Sun is the same at 1 AU as it is at 5 AU then it
can be argued that an ascending CME could still be rooted at the Sun for a
week, or, indeed, much longer.

In reality, an ascending flux rope would most likely contain a mix of open and
closed field lines, driven by apparently random reconnection events (Crooker &
Horbury 2006; Gosling et al. 1995). Complete disconnection of the structure
appears to be unlikely.

With the new STEREO HI data we should be able to test this scenario, and this
has been reported by Harrison et al. (2008). The HI data appear to confirm the
in situ interpretation showing coherent structures, apparently still connected
to the Sun over long distances. There is no evidence for events pinching-off.
However, this in turn presents us with an anomaly. McComas (1995) has argued
that the heliospheric magnetic flux does not continually build up, so flux
must be shed through reconnection somehow during the ICME process. If we are
rejecting the plasmoid or disconnected ICME scenario then we must find another
way of limiting the flux build up over time.

In the absence of evidence for the pinching-off of CMEs, an interchange
reconnection process has been suggested as the mechanism by which CMEs
disconnect from the Sun (Gosling et al. 1995; Crooker et al. 2002). The basic
idea is that the ascending CME can travel a considerable distance, well beyond
the Earth, still connected to the Sun, and that perhaps days or even weeks
after the onset, the legs of the CME, still rooted in the Sun, will interact
with adjacent open field lines at low altitude in the corona; reconnection
results in the formation of low-lying loops as one CME leg reconnects with the
adjacent fields and an outward ascending kink-shaped structure ascends into
the heliosphere from the site of one of the original CME footpoints.

This approach has a few attractive points. For example, it seems logical that
the site of the greatest field density, magnetic complexity and field-line
motion would be the most likely site of any reconnection in the ascending CME.
However, assuming that such interchange reconnection is the ‘end game’ of a
CME, and that this low level reconnection results in the outward propagation
of a kinked field-line configuration, what might we expect to observe and,
indeed, have we seen such features? Harrison et al. (2008) indeed point to
observations of narrow V-shaped structures identified in the HI data that
could be candidates for such reconnection events.

It is early days for this work using STEREO but the indications are that there
is plenty to be gained from these studies. As the mission progresses we
anticipate more opportunities where we have the chance to combine both imaging
and in situ measurements of specific events, and their impacts, as well as to
model CMEs in the heliosphere in 3D as never before. Thus, this report should
be take as an early statement on the progress and direction of this work which
is opening a new chapter in solar, heliospheric and space weather physics.

##  12 Coronal Mass Ejections and Space Weather (N. Gopalswamy)

CMEs cause adverse space weather in two ways: (i) when they arrive at Earth’s
magnetosphere, they can couple to Earth’s magnetic field and cause major
geomagnetic storms (Gosling et al. 1990) and (ii) they can drive fast mode MHD
shocks that accelerate solar energetic particles (Reames 1999). Significant
progress has been made on both these aspects over the past few years. In the
case of geomagnetic storms, connecting the magnetic structure and kinematics
of ICMEs observed at 1 AU to the CME source region at the Sun has received
considerable attention. In the case of SEPs, assessing the contribution from
flare reconnection and shock to the observed SEP intensity has been the focus.
The importance of the variability in the Alfven speed profiles in the outer
corona is also under investigation because of its importance in deciding the
shock formation.

###  12.1 Geomagnetic Storms

High-Speed Solar Wind Streams (HSS) interacting with the slow solar wind
result in corotating interaction regions (CIRs), which also can produce
geomagnetic storms (Vršnak et al. 2007a), but they are generally weaker than
the CME-produced storms (Zhang et al. 2007). Occasionally, the CIR and ICME
structures combine to produce major storms (Dal Lago et al. 2006). Multiple
CMEs are often involved producing some super-intense storms (Gopalswamy et al.
2007; Zhang et al. 2007). There are numerous effects produced by the ICMEs in
the magnetosphere and various other layers down to the ground (see Borovsky et
al. 2006; Kataoka & Pulkkinen 2008).

The key element of ICMEs for the production of geomagnetic storms is the
southward magnetic field component. While the quite heliospheric field has no
out of the ecliptic field component (except for Alfvenic fluctuations in the
solar wind), a CME adds this component to the interplanetary (IP) magnetic
field. If an ICME has a flux rope structure, one can easily see that the
azimuthal component of the flux-rope field or its axial component forms the
out of the ecliptic component. In ICMEs with a flux rope structure (i.e.,
magnetic clouds), it is easy to locate the southward component from the
structure of the cloud (Gopalswamy 2006a; Wang et al. 2007; Gopalswamy et al.
2008a). In non-cloud ICMEs, it is not easy to infer the location of the
southward component. If the ICMEs are shock-driving, then the magnetosheath
between the shock and the driving ICME (Kaymaz & Siscoe 2006; Lepping et al.
2008) can contain southward field and hence cause geomagnetic storms
(Gopalswamy et al. 2008a). The cloud and sheath storms can be substantially
different (Pulkkinen et al. 2007).

Once an IP structure has a southward magnetic field, the efficiency with which
it causes geomagnetic storm depends on the strength of the magnetic field and
the speed with which it hits the magnetosphere (Gonzalez et al. 2007;
Gopalswamy 2008d). Statistical investigations have shown that the storm
intensity (measured e.g., by the Dst index) is best correlated with the speed-
magnetic field product in magnetic clouds and their sheaths. Interestingly, an
equally good correlation is obtained when the magnetic cloud/sheath speed is
replaced by the CME speed measured near the Sun (Gopalswamy et al. 2008a).
This suggests that if one can estimate the magnetic field in CMEs near the
Sun, the strength of the ensuing magnetic storm can be predicted. The ICME
speed can be predicted based on the CME speed by quantifying the interaction
between CMEs and the solar wind (Xie et al. 2006; Nakagawa et al. 2006; Jones
et al. 2007; Vršnak & Zic, 2007). Most of the storm-causing CMEs are halo
CMEs, which are subject to projection effects and hence space speeds cannot be
easily measured (Kim et al. 2007; Gopalswamy & Xie 2008; Howard et al. 2007;
Vršnak et al. 2007b). There have been several attempts to use the sky-plane
speed of CMEs to obtain their space speed (Xie et al. 2006; Michalek et al.
2008; Zhao 2008) with varying extents of success. The magnetic field strength
and kinetic energy of CMEs are somehow related to the free energy available in
the source region. Quantifying this free energy has been a difficult task
(Ugarte-Urra et al. 2007; Schrijver et al. 2008).

The solar sources of CMEs need to be close to the disk center for the CMEs to
make a direct impact on Earth and they have to be fast. In fact the solar
sources of magnetic clouds, storm-causing CMEs, and halo CMEs have been shown
to follow the butterfly diagram suggesting that only sunspot regions have the
ability to produce such energetic CMEs (Gopalswamy 2008d). The average near-
Sun speed of CMEs that cause intense geomagnetic storms is $\sim 1000$ km s-1
(Gopalswamy 2006b; Zhang et al. 2007), similar to the average speed of halo
CMEs (Gopalswamy et al. 2007) because many of the storm-producing CMEs are
halo CMEs. Halo CMEs are more energetic (Lara et al. 2006; Liu 2007;
Gopalswamy et al. 2007, 2008a) and end up being magnetic clouds at 1 AU. Most
halo CMEs ( 70%) are geoeffective. Non-geoeffective halos are generally
slower, originate far from the disk center, and originate predominantly in the
eastern hemisphere of the Sun. The geoeffectiveness rate of halo CMEs has been
reported to be anywhere from 40% to more than 80% (Yemolaev & Yermolaev 2006),
but the difference seems to be due to different definitions used for halo CMEs
(some authors have included all CMEs with width $>120^{\rm o}$ as halos) and
the sample size (Gopalswamy et al. 2007). The geoeffectiveness rate of CMEs
may be related to the fact that more ICMEs are observed as magnetic clouds
during solar minimum than during solar maximum (Riley et al. 2006). It is
possible that all ICMEs are magnetic clouds if viewed appropriately (Krall
2007). This suggestion is consistent with the ubiquitous nature of post
eruption arcades, which seem to indicate flux rope formation in the eruption
process (Kang et al. 2006; Qiu et al. 2007; Yurchyshyn 2008). While the
reconnection process certainly forms a flux rope, it is not clear if the
reconnection creates a new flux rope or fattens an existing one.

###  12.2 SEP Events

Energetic storm particle (ESP) events are the strongest evidence for SEP
acceleration by shocks, but this happens when the shocks arrive at the
observing spacecraft near Earth (Cohen et al. 2006). This means the shocks
must have been stronger near the Sun accelerating particles to much higher
energies. The strongest evidence for SEPs in flares is the gamma-ray lines,
which are now imaged by RHESSI (Lin 2007). All shock-producing CMEs are
associated with major flares (M- or X-class in soft X-rays), so both
mechanisms must operate in most SEP events. There has been an ongoing debate
as to which process is dominant based on SEP properties such as the spectral
and compositional variability at high energies (Tylka & Lee 2006; Cane et al.
2007).

The easiest way to identify shocks near the Sun are the type II radio bursts
especially at frequencies below 14 MHz, which correspond to the near-Sun IP
medium (Gopalswamy 2006c). Analyzing electrons and protons in SEP events,
Cliver & Ling (2007) have found evidence for a dominant shock process
including flatter SEP spectra, apparent widespread sources, and high
association with long wavelength type II bursts. A recent statistical study
finds the SEP association rate of CME steadily increases with CME speed and
width especially and there is one-to-one correspondence between SEP events and
CMEs from the western hemisphere with long wavelength type II bursts
(Gopalswamy et al. 2008b). Type II burst studies have also have concluded that
the variability in Alfven speed in the outer corona decides the formation and
strength of shocks (Shen et al. 2007; Gopalswamy et al. 2008c). For example, a
400 km s-1 CME can drive a shock, while a 1000 km s-1 CME may not drive a
shock, depending on the local Alfven speed.

James A. Klimchuk

President of the Commission

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