Chromospheric Flares
Coimbra Solar Physics Meeting on The Physics of Chromospheric Plasmas
ASP Conference Series, Vol. xxx, 2007
Petr Heinzel, Ivan Dorotovič and Robert J. Rutten, eds.
Chromospheric Flares
Hugh S. Hudson
Space Sciences Laboratory, University of California, Berkeley
Abstract. In this topical review I revisit the “chromospheric flare”. This
should currently be an outdated concept, because modern data seem to rule out
the possiblity of a major flare happening independently in the chromosphere
alone, but the chromosphere still plays a major observational role in many ways.
It is the source of the bulk of a flare’s radiant energy – in particular the visi-
ble/UV continuum radiation. It also provides tracers that guide us to the coronal
source of the energy, even though we do not yet understand the propagation of
the energy from its storage in the corona to its release in the chromosphere. The
formation of chromospheric radiations during a flare presents several difficult
and interesting physical problems.
1. Introduction
Solar flares first revealed themselves as visual perturbations of the solar atmo-
sphere (“white light flares”) and hence immediately were construed as a pho-
tospheric process. With the invention of spectroscopic techniques, though, it
became clear that chromospheric emission lines such as Hα revealed flare pres-
ence much more readily. This led to the concept of the “chromospheric flare” and
to a great deal of observational material on Hα flares and eruptions, as reviewed
by Smith & Smith (1963), Zirin (1966), or Švestka (1976), for example. At some
point, prior to the discovery of coronal flare effects, the misinterpretations of the
Hα line profile even led to the incorrect idea that a flare was a sudden cooling of
the solar atmosphere. In any case, a perturbation of the lower solar atmosphere
violent enough to affect the solar luminosity itself (“white light”) implies a large
energy content.
Our view of flares now emphasizes the high temperatures and non-thermal
effects seen in the corona, and we generally believe the chromospheric effects
themselves to be secondary in nature. This may be true, but nonetheless the
modern observations confirm the fact that the lower solar atmosphere dominates
the radiant energy budget of a flare via the UV and white-light continua. Some-
how, therefore, the energy stored in the solar corona rapidly focuses down into
regions visible in chromospheric signatures; this accounts for the high contrast
of flare effects there. Thus the “chromospheric flare” remains essential to our
understanding of the overall processes involved.
The chromosphere nowhere exists as a well-defined layer with a reproducible
height structure. In this paper I use the term interchangeably with “lower
solar atmosphere,” embracing the phenomena of the visible photosphere through
the transition region. During flares the structure of these “layers” and the
physical conditions within them may change drastically. The changes generally
http://arxiv.org/abs/0704.0823v1
2 Hudson
happen so fast and on such small spatial scales that we cannot observe them
comprehensively. Understanding the impulsive phase in the chromosphere may
therefore seem like something of a lost cause from the the point of view of
theory, especially in view of our inability to understand the quiet chromosphere
any better than we do. The data repeatedly reveal that we simply have not
yet resolved the spatial or temporal structures involved in the impulsive phase,
and that without knowing the geometry of the physical structure, we cannot
really comprehend its physics. The TRACE (Handy et al. 1999) and RHESSI
(Lin et al. 2002) observations have provided more than one recent breakthrough,
however, and it may be that we are beginning to understand the gradual phase
of a flare at least.
This review is organized around several topics involving the behavior of
the chromosphere during a flare. These include the process of “chromospheric
evaporation” (Section 4), flare energetics (Section 5), the mechanisms of flare
continuum emission (Section 6), and the inference of flare structure from the
morphology of the chromospheric flare (Section 7 and Section 9). In Section 2
and Section 3 we give an overview of the history of chromospheric flares and
show a cartoon to establish a working model of a solar flare. Sections 8 and 11
discuss large-scale magnetic reconnection and theoretical ideas, and Section 10
presents a γ-ray mystery.
2. Historical Development
Although it was the white-light continuum that initially revealed the existence
of solar flares, the advent of spectroscopy (e.g., Hale 1930) allowed their regular
observation via the Hα line (see Švestka 1966 for a discussion of the historical de-
velopment of these observations). This strong absorption line actually becomes
an emission line during bright flares, and Hα limb observations frequently show
prominences and eruptions. Hα observers came to recognize a particular flare
morphology, the so-called two-ribbon flare. Bruzek (1964) described the pat-
terns followed by these events, which provided strong evidence that the solar
corona had to play a major role in flare development. Figure 1 reproduces one
of Bruzek’s sketches, and then illustrates in a cartoon (due to Anzer & Pneuman
1982) how this morphology led to our standard magnetic-reconnection scenario
that tries to embrace the X-ray observations and the coronal mass ejections
(CMEs) as well as the chromospheric ribbon structures.
In this standard picture a solar flare develops in a complicated manner that
involves restructuring of the coronal magnetic field in such a way as to release
energy. The immediate effects of this energy release are to produce broad-band
“impulsive phase” emissions and to drive chromospheric gas up into coronal
magnetic loops, the process we term “chromospheric evaporation.” A part of the
field magnetic structure may actually erupt and open out into the solar wind, in
the sense that the field lines stretch out past the Alfvé n critical point of the flow.
This opening may consist of rising loops which then take the form of a coronal
mass ejection (CME), or it may involve interactions with previously open field (a
process often termed “interchange reconnection” nowadays; see Heyvaerts et al.
1977). If a CME does accompany the flare, as it almost invariably does for
flares of GOES class X or greater, the energy involved in mass motions may
Chromospheric Flares 3
Figure 1. Left: one of Bruzek’s (1964) sketches, showing a flare with ribbons
on the disk and its equivalent Hα “loop prominence system” over the limb.
This key observational pattern led directly to the formation of our standard
flare model (right), in the form presented by Anzer & Pneuman (1982).
be comparable to the luminous energy (e.g., Emslie et al. 2005). Generally the
observations are limited in resolution, both temporal and spatial, and especially
in spectral coverage. Thus we often resort to a cartoon that serves to identify
how the essential parts of a flare relate to one another.
Soft X-ray observations show hot loops in the gradual phase of a flare.
These result from the material “evaporated” from the chromosphere and have
anomalously high gas pressure (but still low plasma β; however see Gary 2001).
Whereas the pressure at the base of the corona normally is of order 0.1 dyne cm−2,
a bright flare loop can achieve 103 dyne cm−2. This over-dense and over-hot coro-
nal loop gradually cools, and in its final stages the remaining plasma returns
to a more chromospheric state and suddenly becomes visible in Hα (Goldsmith
1971). The loops that have reached this state then form Bruzek’s Hα loop
prominence system (Figure 1).
During the ribbon expansion another important phenomenon occurs: hard
X-ray emission appears at the footpoints of the coronal loops that are in the pro-
cess of being filled by chromospheric evaporation (Hoyng et al. 1981). The hard
X-rays show that a substantial part of flare energy appears in the form of non-
thermal electrons (Kane & Donnelly 1971; Lin & Hudson 1976; Holman et al.
2003). The hard X-ray signature (and hence the energetic dominance of these
electrons) is present whether or not the flare develops the two-ribbon morphol-
ogy or has a CME association.
The hard X-ray emission occurs in the impulsive phase of the flare, contem-
poraneously with the period of chromospheric evaporation that fills the coro-
nal loops and with the acceleration phase of the associated CME (Zhang et al.
2001). In Section 5 we describe this phase of the flare with the thick-target
model (Kane & Donnelly 1971) which Hudson (1972) identified with the energy
source of white-light flare continuum.
3. The Flare Spectrum
A (major) flare can be observed at almost any wavelength in a fast-rise/slow-
decay time profile, with some (e.g., the white-light continuum) having a more
impulsive variability, and others (e.g., the Balmer lines) having a more gradual
pattern (Figure 2, right). We generally describe a flare as consisting of a foot-
4 Hudson
Figure 2. Left: Line widths of the Balmer-series lines, from the classic
paper by Suemoto & Hiei (1959). The inferred densities added to the curves
are logne = 13.5 and 13.3; the inferred filling factor is small, suggesting
either filamentary structure or thin layering. Right: Typical time series of
flare radiations, distinguishing the impulsive phase from the gradual phase
(see Kane & Donnelly 1971).
point and ribbon structures in the lower atmosphere, coronal loops, and various
kinds of ejecta. The impulsive phase is typically associated with the footpoint
structures, and the gradual phase with the flare ribbons. Nowadays imaging
spectroscopy in principle allows us to study these regions independently.
Flare spectroscopy began with the observation of the Balmer series, which
shows broad lines tending towards emission profiles as the flare gets more en-
ergetic. Early observations of the higher members of the sequence allowed the
inference of a relatively high density and of a small filling factor (Suemoto & Hiei
1959; see the left panel of Figure 2). Such observations refer to what we would
now call the gradual phase of the flare (see the right panel of Figure 2 for a
sketch of the temporal development of a flare). In the impulsive phase the con-
tinuum appears in emission, as noted originally by Carrington and by Hodgson
independently. The weak photospheric metallic lines may also go weakly into
emission (or are filled in by continuum) and the recent observations of Xu et al.
(2004) show that flare effects can appear even at the “opacity minimum” region
of the spectrum, where one would expect much higher densities. In fact a single
density could never properly describe such a heterogeneous structure, but each
spectral band provides its own clues. At the time of writing no proper analy-
sis of spectroscopic “response functions” (e.g., Uitenbroek 2005) for any of the
signatures has yet been attempted, so our inference of flare structure from the
spectroscopy alone is weak.
The continuum radiation seen in white light and the UV constitutes the bulk
of flare radiated energy (Kane & Donnelly 1971; Woods et al. 2006). TRACE
imaging of this emission component shows it to consist of unresolved, intensely
Chromospheric Flares 5
bright fine structures (Hudson et al. 2006). The thick-target model invokes fast
electrons (energies above about 10 keV) to transport coronal energy into the
chromosphere. Here collisional losses provide the heating and footpoint emis-
sions that accompany the hard X-ray bremsstrahlung. The thick-target model
does not explain the particle acceleration, nor show how the footpoint sources
can be so intermittent. We return to this question in Section 7.
The spectra emitted at the footpoints of the flaring coronal loops have
contributions over an exceptionally broad wavelength range, as sketched in the
right panel of Figure 2. The prototypical observable is the hard X-ray flux, which
imaging observations show to be concentrated at the footpoints (Hoyng et al.
1981), but impulsive footpoint emissions also occur in many spectral windows
ranging from the microwaves (limited presumably by opacity) to the γ-rays
(limited presumably by detection sensitivity). There is a large body of work
on the Hα line alone, both observation and theory. Berlicki (2007) reviews the
Hα spectroscopic material in detail in these proceedings. A strong absorption
line forms across a wide range of continuum optical depths, and in principle this
single line might provide sufficient information to infer the physical structure of
the flare everywhere. In practice the complexities of the radiative transfer and
of the flare motions, especially in the impulsive phase, make this information
ambiguous (see Berlicki 2007).
4. Chromospheric Evaporation
The motions most directly relevant to the chromosphere are often called “chro-
mospheric evaporation,” even though the direct Doppler signatures of this mo-
tion are normally found in lines formed at higher temperatures (but see Berlicki et al.
2005). That this process occurs (even if it is not “evaporation” strictly speak-
ing) was suggested by the early observations of loop prominence systems (e.g.,
Bruzek 1964) with their “coronal rain,” and Neupert (1968) established its as-
sociation with non-thermal processes such as bursts of microwave synchrotron
radiation. The thermal microwave spectrum (e.g., Hudson & Ohki 1972) made
it particularly clear that the gradual phase of a solar flare involves the temporary
levitation of chromospheric material into the corona, as opposed to the process
that might be imagined from the earlier term “sporadic coronal condensation”
(e.g., Waldmeier 1963). The flows involved in chromospheric evaporation are
along the field direction and serve to create systems of coronal loops with rel-
atively high gas pressure and therefore higher (but still probably low) plasma
beta.
The early observational indications of chromospheric evaporation actually
came from blueshifts in EUV and soft X-ray lines (e.g., Antonucci et al. 1982;
Acton et al. 1982) such as those from FeXXV or CaXIX. Figure 3 shows an
image-resolved view of Doppler shifts in an evaporative flow (Czaykowska et al.
1999). The chromospheric effects are more subtle and in fact the impulsive-phase
evaporation is difficult to disentangle from other effects (Schmieder et al. 1987).
The high-temperature blueshifts correspond to upward velocities of some hun-
dreds km/s and seldom appear in the absence of a stationary emission line; in
other words, hot plasma has already accumulated in coronal loops as the process
continues. Based on theory and simulations (Fisher et al. 1985) one can distin-
6 Hudson
Figure 3. Imaging spectroscopy from SOHO/CDS of EUV emission lines
in the gradual phase of a two-ribbon flare, showing the clear signature of
blueshifted upflows in the expected locations along the flare ribbons. This is
“gentle” evaporation not associated with strong hard X-ray emission (from
Czaykowska et al. 1999). Note that CDS produces images by scanning in one
spatial dimension, so that each image (while monochromatic) is not instanta-
neous.
guish “explosive” and “gentle” evaporation, depending upon the physics of en-
ergy deposition (e.g., Abbett & Hawley 1999). In explosive evaporation, driven
hypothetically by an electron beam, one has the additional complication of a
“chromospheric condensation” that produces a redshift as well. Schmieder et al.
(1987) and Berlicki et al. (2005) survey our overall understanding. It would be
fair to comment that the explosive evaporation stage remains ill-understood,
even though it principle it describes the key physics of sudden mass injections
into flare loops.
Chromospheric Flares 7
-500 0 500 1000 1500 2000 2500 3000
VAL C Height, km
-500 0 500 1000 1500 2000 2500 3000
VAL C Height, km
10000
Figure 4. Characteristic radiative cooling time (upper) as a function of
height in the VAL-C model, crudely estimated as described in the text. The
lower panel shows the temperature in this model.
5. Energetics and Magnetic Field
We can use the standard VAL-C model (Vernazza et al. 1981), as discussed
further in the Appendix, to discuss the energetics. First we establish that the
chromosphere and the rest of the lower solar atmosphere (i.e., that for which
τ5000 < 1) have negligible heat capacity and limited time scales. Figure 4 shows
an estimate of the radiative time scale in the VAL-C model (Vernazza et al.
1981). This shows 3σ(z)kT/L⊙, where σ is the surface density as a function of
height about the τ5000 = 1 layer, and L⊙ the solar luminosity. The time scale
decreases below 1 sec only above z ∼ 515 km, near the temperature minimum
in the VAL-C model. Above this height any energy injected into the system will
tend to radiate rapidly, resulting in a direct energy balance between input and
output energy, rather than a local storage and release. At lower altitudes we
would not expect to see rapid variability.
The model also allows us to ask whether the chromosphere itself can store
energy comparable to that released in a major flare or CME. Table 1 gives some
order-of-magnitude properties for a chromospheric area of 1019 cm2, showing
both possible sources (bold) and sinks (italics) of energy. For the magnetic field
we simply assume 10 or 1000 G as representative cases. Using the total magnetic
energy in this manner is an upper limit, since the actual free energy would
depend on its degree of non-potentiality. We find that magnetic energy storage
limited to the volume of the chromosphere will not suffice, unless unobservably
small-scale fields there somehow dominate. The gravitational potential energy
also will not suffice. Estimates of this sort confirm the idea that the flare energy
must reside in the corona prior to the event.
The estimate of gravitational potential energy is somewhat more ambiguous.
The Table shows the value needed to displace the entire atmosphere by its total
thickness, the equivalent of roughly 3′′ in the VAL-C model. There does not
seem to be any evidence for such a displacement, although I am not aware of
any searches. It is likely that the stresses that store energy in the coronal field
8 Hudson
have their origin deeper in the convection zone, rather than in the atmosphere
(McClymont & Fisher 1989). Actually the observable changes of gravitational
energy are even of the wrong sign, given that we normally observe only outward
motions, (against gravity) during a flare.
Table 1. Properties of a chromospheric volume of area 1019 cm2
Parameter VAL-C VAL-C above Tmin
Mass 4×1019 gram 5×1017 gram
Magnetic energy 1×1028erg 8×1027 erg (10 G)
Magnetic energy 1×1032 erg 8×1031 erg (103 G)
Gravitational energya 3×1032 erg 3×1030 erg
Thermal energy 2×1031 erg 3×1029 erg
Kinetic energy 3×1029 erg 3×1027 erg
Ionization energy 4×1032 erg 6×1030 erg
aPotential energy for a vertical displacement of 2.5 × 108 cm
From Table 1 one concludes that the chromosphere probably does play a
dominant role in the energetics of a solar flare, at least as described by a semi-
empirical model such as VAL-C. This just restates the conventional wisdom,
namely that the flare energy needs to be stored magnetically in the corona,
rather than in the chromosphere where the radiation forms. Note that this
is backwards from the relationship for steady emissions: the requirement for
chromospheric heating is larger than that for coronal heating, so it is possible to
argue that the steady-state corona actually forms as a result of energy leakage
from the process of chromospheric heating (e.g., Scudder 1994).
We can make a similar estimate for energy flowing up from the photosphere.
The Alfvén speed at τ5000 = 1 ranges from 3 to 30 km s
−1, depending on the
field strength, in the VAL-C model (see Appendix). Below the surface of the Sun
vA drops rapidly because of the increase of hydrogen ionization. Thus chromo-
spheric flare energy cannot have been stored just below the photosphere, since it
could not propagate upwards rapidly enough (McClymont & Fisher 1989). This
again supports the conventional wisdom that flare energy resides in the corona
prior to the event.
To drive chromospheric radiations from coronal energy sources requires ef-
ficient energy transport, which is normally thought to be in the form of non-
thermal particles (the “thick target model”, Brown 1971; Hudson 1972) or in
the form of thermal conduction as in the formation of the classical transition
region. Both of these mechanisms provide interesting physical problems, but the
impulsive phase of the flare (where the thick-target model usually is thought to
apply) certainly remains less understood. Section 11 comments on models.
The magnetic field in the chromosphere is decisively important but ill-
understood. The plasma beta is generally low (see Appendix), so just as in
the corona the dynamics depends more on the behavior of the field itself than
to the other forces at work. Generally we believe that the subphotospheric field
exists in fibrils, implying the existence of sheath currents that isolate the flux
tubes from their unmagnetized environment. On the other hand, the dominance
of plasma pressure in the chromosphere as well as the corona implies that the
Chromospheric Flares 9
field must rapidly expand to become space-filling. Longcope & Welsch (2000)
discuss the physics involved in this process as flux emerges from the interior.
The effect of the flux emergence must be to create current systems linking the
sources of magnetic stress below the photosphere, with the non-potential fields
containing the coronal free energy. A full theory of how this works does not
exist, and we must add to the uncertainty the possibilty of unresolved fields
(e.g., Trujillo Bueno et al. 2004). Their suggested factor of 100 in B2 would
clearly affect the estimate of magnetic energy given in Table 1 and perhaps
change everything. We note in this context that the “impulse response” flares
(White et al. 1992) have scales so small that one could argue for an entirely
chromospheric origin.
6. Energetics and the Formation of the Continuum
The formation of the optical/UV emission spectrum of a solar flare has from
the outset presented a special challenge, since (a) it represents so much energy,
and (b) it appears in what should be the stablest layer of the solar atmosphere.
The recent observations of rapid variability and spatial intermittency make this
all the more interesting, and these observations – now from space – also help
to intercompare events; previous catalogs of white-light flares (e.g., Neidig 1989
and references therein) had to be based on spotty observations made with a
wide variety of instruments. Observationally, the continuum appears to have two
classes, with most events (“Type I” spectra) showing evidence for recombination
radiation via the presence of the Balmer edge and sometimes the Paschen edge
as well. A few events (e.g., Machado & Rust 1974) show spectra with weak or
unobservable Balmer jumps, implicating H− continuum as observed in normal
photospheric radiation. The spectra in the latter class (“Type II”) suggests
a relationship to Ellerman bombs (Chen et al. 2001). However, the physics of
Ellerman bombs appears to be quite different from that of solar flares (e.g.,
Pariat et al. 2004), though.
The strong suggestion from correlations is that non-thermal electrons phys-
ically transport flare energy from the corona, where it had been stored in the
current systems of non-potential field structures, into the radiating layers. The
hard X-ray bremsstrahlung results from the collisional energy losses of these
particles, and other signatures (such as the optical/UV continuum) depends on
secondary effects. Proposed mechanisms include direct heating, heating in the
presence of non-thermal ionization, and radiative backwarming. In some manner
these effects (or others not imagined) must provide the emissivity ǫν , to support
the observed spectrum. Note that the emissivity is often expressed in terms
of the source function Sν = ǫν/κν via the opacity κν . In a steady state one
would have energy balance between the input (e.g., electrons) and the contin-
uum. Fletcher et al. (2007) have now shown that this implies energy transport
by low-energy electrons, below 25 keV, as opposed to the 50 keV or higher sug-
gested by some earlier authors. Such low-energy electrons have little penetrating
power and could not directly heat the photosphere itself from a coronal accel-
eration site. Thus either the continuum arises from altered conditions in the
chromosphere, or some mechanism must be devised to link the chromosphere
and photosphere not involving the thick-target electrons.
10 Hudson
“Radiative backwarming” – for example Balmer and Paschen continuum
excited in the chromosphere and then penetrating down to and heating a deeper
layer – could in principle provide a vertical step between energy source and sink.
One problem with this is that the weaker backwarming energy fluxes might
not cause appreciable heating in the denser atmosphere, and thus not be able
to contribute to the observed continuum excess, because of the short radiative
time scale. This idea is a variant of the mechanism of non-thermal ionization
originally proposed by Hudson (1972) in the “specific ionization approximation,”
which involves no radiative-transfer theory and simply assumes ion-electron pairs
to be created locally at a mean energy (∼30 eV per ion pair). Finally, the rapid
variability observed in the continuuum, even at 1.56 µm (Xu et al. 2006) provides
a clear argument that the continuum forms at the temperature minimum or
above (see Section 5, especially Figure 4).
Early proponents of particle heating as an explanation for white-light flares
also considered protons as an energy source (Najita & Orrall 1970; Švestka
1970). This made sense, because protons at energies even below those char-
acteristic of γ-ray emission-line excitation can penetrate more deeply than the
electrons that produce hard X rays. It makes even more sense now that we have
the suggestion that ion acceleration in solar flares may rival electron accelera-
tion energetically (Ramaty et al. 1995). Simnett & Haines (1990) suggest that
particle acceleration in solar flares involves a neutral beam, implying that the
major energy content (and hence the optical/UV continuum) would originate in
the ion component. This idea does not appear to explain the apparent simul-
taneity of the footpoint sources (Sakao et al. 1996), and at present we do not
understand the plasma physics of the particle acceleration and propagation well
enough even to identify the location of the acceleration region.
7. Flare Structures Inferred from Chromospheric Signatures
The continuum kernels may move systematically for perhaps tens of seconds and
generally have short lifetimes. We illustrate this in Figure 5 (from Fletcher et al.
2004). This shows the motions of individual UV bright points within the flare
ribbon structure. Such motions are only apparent motions, as in a deflagration
wave, because they exceed the estimated photospheric Alfvén speed (see Sec-
tion 4 and the Appendix). Figure 6 (from Hudson et al. 2006) makes the same
point for a different flare, using TRACE white-light observations. The basic
picture one gets from such observations is that the white light/UV continuum
of a flare appears in compact structures that are essentially unresolved in space
and in time within the present observational limits. These bright points con-
tain enormous energy and thus must map directly to the energy source. We do
not know if the fragmentation (intermittency) results from this mapping or is
intrinsic to the basic energy-release mechanism.
How do the small chromospheric sources map into the corona, where the
flare energy must reside on a large scale before its release? A strong literature
has grown up regarding this point, interpreting the ribbon motions as measures
of flux transfer in the standard magnetic-reconnection model (Poletto & Kopp
1986; see also literature cited, for example, by Isobe et al. 2005). The flux
transfer in the photosphere is taken to measure the coronal inflow into the re-
Chromospheric Flares 11
Figure 5. Flare footpoint apparent motions deduced from TRACE UV ob-
servations. Each squiggle represents the track of a bright point visible for
several consecutive images at a few-second cadence, with the black dot show-
ing the beginning of each track (Fletcher et al. 2004).
Figure 6. Intermittent structure seen in TRACE white-light images of
an M-class flare on July 24, 2004. The individual frames have dimensions
32′′× 64′′. Note the presence of bright features consistent with the TRACE
angular resolution, and which change from frame to frame over the 30-second
interval. These observations do not appear to resolve the fluctuations either
in space or in time (Hudson et al. 2006).
connecting current sheet, which appears to correlate with the radiated energy as
seen in hard X-rays, UV, or Hα. Figure 1 (right) shows the assumed geometry
linking the chromosphere and corona. The analysis extends to the multiple si-
multaneous UV footpoints apparently moving within the ribbons as they evolve,
as noted in Figure 5 above. The analyses suggest a strong relationship between
energy release and the inferred coronal Alfvén speed.
12 Hudson
Figure 7. Left: UV ribbons (TRACE observations) from a flare of Novem-
ber 23, 2000. The gray scale shows the time sequence of brightening. Right:
Correlation between pixel brightness in Ribbon A and the inferred reconnec-
tion rate (from Saba et al. 2006).
8. Dynamics and Magnetic Reconnection
To release energy from coronal magnetic field in a largely “frozen-field” plasma,
a flare must involve mass motions. We often do observe apparent motions,
both parallel and perpendicular to the field as indicated by the image striations
(“loops”). Most of the observable motions are outward, leading to the idea of a
“magnetic explosion” (e.g., Moore et al. 2001). Motions apparently perpendicu-
lar to the magnetic field may become coronal mass ejections (CMEs) and contain
a great deal of energy (e.g., Emslie et al. 2005). These perpendicular motions
also are involved in flare energy release; for example the large-scale magnetic
reconnection involved in many flare models (Figure 1, right panel) necessarily
involve “shrinkage” (e.g., Švestka et al. 1987; Forbes & Acton 1996). Note that
this process is more of a magnetic implosion than a magnetic explosion (Hudson
2000).
The motion of flare footpoints and ribbons is (we believe) only apparent,
because of the low Alfvén speed vA in the photosphere, where the field is tem-
porarily anchored (“line-tied”). For B =1000 G and n = 1017 cm−3 we find
vA ∼ 6 km s
−1; observations often suggest motions an order of magnitude faster
(e.g., Schrijver et al. 2006). The motions therefore represent a wave-light confla-
gration moving through a relatively fixed magnetic-field structure. It is natural
to imagine that this sequence of field lines links to the coronal energy-release
site, which the standard model identifies with a current sheet that mediates
large-scale magnetic reconnection.
Figure 7 shows one example of the result of an analysis of the apparent
motion of a flare ribbon (Saba et al. 2006). This and other similar analyses reveal
a tendency for the “reconnection rate” to correlate with the pixel brightness. The
reconnection rate is the rate at which flux is swept out in the ribbonmotion, often
expressed as an electric field from E = v×B (the so-called “reconnection electric
field”). In this picture the flare ribbons are identified with “quasi-separatrix
structures” where magnetic reconnection can take place most directly.
Chromospheric Flares 13
9. Surges, Sprays, and Jets
Chromospheric material also appears in the corona in the form of surges and
sprays, which may have a close relationship to the flare process (e.g., Engvold
1980). In addition, of course, we observe filaments and prominences in chromo-
spheric lines, and these also have a flare/CME association, but too tangential
for discussion in this review.
Surges and sprays are Hα ejecta, rising into the corona as a result of chro-
mospheric magnetic activity. The literature traditionally distinguishes them by
apparent velocity, with the faster-moving sprays taken to have stronger flare
associations. Surges often appear to return to the Sun, while sprays acceler-
ate beyond the escape velocity and do not return. Both appear to move along
the magnetic field lines, but unlike the evaporation flow the surges and sprays
incorporate material at chromospheric temperatures.
Modern soft X-ray and EUV data (Yohkoh, SOHO, and TRACE) have had
sufficient time resolution to reveal the phenomenon of X-ray jets (Shibata et al.
1992); see also the UV observations of Dere et al. 1983. These tightly-collimated
structures at X-ray temperatures have a strong correlation with surges and
sprays, and indeed presumably lead to the jet-like CMEs seen at much greater
altitudes (Wang & Sheeley 2002). These events have a strong association with
emerging flux, and indeed the X-ray jets invariably have an association with mi-
croflares and originate in the chromosphere near the microflare loop(s) (Shibata et al.
1992). As Zirin famously remarked, most emerging flux emerges within active
regions, and that is where the jets preferentially occur. The site is frequently
in the leading part of the sunspot group. Figure 8 (Canfield et al. 1996) shows
the sequence of events in an explanation of these phenomena invoking magnetic
reconnection to allow chromospheric material access to open fields. Note that
this scenario imposes two requirements on the chromosphere: there must be
open and closed fields juxtaposed, and a large-scale reconnection process must
be able to proceed under chromospheric conditions. The Canfield et al. (1996)
observations strongly imply that this process requires the presence of vertical
electric currents supporting the observed twisting motions.
The surges, sprays, and jets, not to mention flares and CMEs, underscore
the time dependence and three-dimensionality characterizing what is often char-
acterized as a thin time-independent layer for convenience. The subject of
spicules is outside the scope of this review, but we note that they represent
a form of activity that occurs ubiquitously outside the magnetic active regions.
10. A Chromospheric γ-ray Mystery
The γ-ray observations of solar flares have begun, as did the radio and X-ray
observations before them, to open new windows on flare physics. Share et al.
(2004) have made a discovery that is difficult to understand and which in-
volves chromospheric material. They report observations of the line width of the
0.511 MeV γ-ray emission line formed by positron annihilation (Figure 9). This
emission requires a complicated chain of events: the acceleration of high-energy
ions, their collisional braking and nuclear interactions in the solar atmosphere,
the emission of secondary positrons by the excited nuclei, the collisional braking
14 Hudson
Figure 8. A mechanism for jet/surge formation involving emerging flux
(upper left), with magnetic reconnection against already-open fields (upper
right), which may lead to a high-temperature ejection (the jet) entraining
chromospheric material (the surge). The cartoon at lower right describes the
observations of (Canfield et al. 1996), who find a spinning motion suggesting
that the process must occur in a 3D configuration rather than that of the
cartoons left and above.
of these energetic positrons in turn, and finally their recombination with ambient
electrons to produce the 0.511 MeV γ-rays. Because the γ-ray observations are
so insensitive, this process requires an energetically significant level of particle
acceleration that is possibly distinct from the well-known electron acceleration
in the impulsive phase.
The mystery comes in the line width of the emission line. Surprisingly the
pioneering RHESSI observations of Share et al. (2004) showed it to be broad
enough to resolve. The likeliest source of this line broadening is Doppler mo-
tions in the positron-annihilation region. This requires the existence of a large
column density (of order gram cm−2) at transition-region temperatures; the
transition region under hydrostatic conditions would be many orders of mag-
nitude thinner (see also Figure 11). According to Schrijver et al. (2006), the
excitation of the footpoint regions during the the time of intense particle accel-
eration only continues for some tens of seconds at most. This would represent
the time scale for the apparent motion of a foopoint source across its diameter.
The γ-ray observations, on the other hand, require minutes of integration for a
statistically significant line-profile measurement.
We therefore are confronted with a major problem. What is the structure
of the flaring atmosphere that permits the formation of the broad 0.511 MeV
γ-ray line? Recent spectroscopic observations of the impulsive phase in the UV,
as viewed off the limb (Raymond et al. 2007) make a conventional explanation
difficult.
Chromospheric Flares 15
500 505 510 515 520
Energy (keV)
0.000
0.005
0.010
0.015
Figure 9. RHESSI γ-ray observations of the 0.511 MeV line of positron
annihilation (Share et al. 2004). The two line profiles are from different inte-
grations in the late phase of the X17 flare of 28 October 2003; for the broader
line the authors suggest thermal broadening, which would require a large
column depth of transition-region temperatures during the flare.
11. Theory and Modeling
To understand the chromospheric spectrum of a solar flare we must understand
the formation of the radiation and its transfer in the context of the motions
produced by (or producing) the flare. The representation of the spectrum by a
“semi-empirical model” represents one shortcut; in such an approach (e.g., the
standard VAL model that we use in the Appendix) one attempts to construct
a model atmosphere capable of describing the spectrum even if it may not be
physically self-consistent. Such descriptions may however be sufficient in the
gradual phase of a flare when the flare loops no longer have energy input and
simply evolve by cooling and draining. Even here, however, we do not have a
good understanding of the “moss” regions that form at the footpoints of these
high-pressure loops (but see Berger et al. 1999). So far as I am aware there is no
literature specifically on “spreading moss,” the similar structure that appears in
association with flare ribbons.
A more complete approach to the physics comes from “radiation hydrody-
namics” physical models, most recently those of Allred et al. (2005); see Berlicki
(2007) for a fuller description. Such models solve the equations of hydrodynam-
ics and radiative transfer simultaneously and can thus deal with chromospheric
evaporation and the formation of the high-pressure flare loops. This frame-
work is necessary if we are to be able to understand the flare impulsive phase
(e.g., Heinzel 2003). Even these models do not have sufficient realism, though,
since they work currently in one dimension and thus cannot follow the time
development of the excitation properly; the high-resolution observations of UV
and white light by TRACE clearly show that the energy release has unresolved
scales. Further, as pointed out by Hudson (1972), the ionization of the chromo-
sphere (and hence the formation of the continua) cannot be described by a fluid
16 Hudson
Figure 10. Continuum emission in the near infrared (1.56 µm, the “opacity
minimum” region) during an X10 flare (Xu et al. 2004). Red shows the IR
emissions, contours show the RHESSI 50-100 keV X-ray sources. The IR
contrast relative to the preflare photosphere reached ∼20% in this event.
approximation, or even by non-LTE radiative transfer that assumes a unique
temperature.
At present there has been little effort to create an electrodynamic theory
of chromospheric flare processes, even though non-thermal particles are widely
thought to provide the dominant energy in at least the impulsive phase. In
the gradual phase there is interesting physics associated with heat conduction
because the transition region would have to become so steep that classical con-
ductivity estimates have difficulty (Shoub 1983). A more complete theory would
have to take plasma effects into account and would probably contain elements
of theories of the terrestrial aurora that are now largely missing from the solar
lexicon. This lack of self-consistency in the modeling probably means that we
have major gaps in our understanding of, for example, the evaporation process
as it affects the fractionation of the elements and of the ionization states of the
flare plasma. The Appendix gives estimates of the ranges of some the key plasma
parameters in the chromosphere.
12. Conclusions
This article has reviewed chromospheric flare observations from the point of view
of the newest available information – Yohkoh, SOHO, TRACE and RHESSI, for
example, but not Hinode or STEREO (already launched), nor much less FASR
or ATST (not launched yet at the time of writing). spite of the high quality of
the data prior to these missions, we still find major unsolved problems:
• How does the chromosphere obtain all of the energy that it radiates?
• How can flare effects appear at great depths in the photosphere?
• How is the anomalous 0.511 MeV line width produced?
Chromospheric Flares 17
• What are the elements of an electrodynamic theory of chromospheric flares?
In my view the solution of these problems cannot be found in chromospheric
observations alone, because the physical processes involve much broader regions
of the solar atmosphere. Even providing answers to these specific questions may
not reveal the plasma physics responsible for flare occurrence, which may involve
spatial scales too fine ever to resolve. But we can hope that new observations
from space and from the ground, in wavelengths ranging from the radio to the
γ-rays, will enable us to continue our current rapid progress, and can speculate
that eventually numerical tools will supplement the theory well enough for us to
achieve full comprehension of the important properties of flares. To get to this
point we will need to deal with the chromosphere, as messy as it is.
One important task that is probably within our grasp now is the compu-
tation of response functions for physical models of flares. At present these are
restricted to very limited numerical explorations of the radiative transfer within
the framework of one-dimensional radiation hydrodynamics (e.g., Allred et al.
2005). The energy transport in these models has been restricted to simplistic
representations of particle beams for energy transport, and do not take account
of complicated flare geometries, waves, or various elements of plasma physics.
Future developments of chromospheric flare theory will need to complete the
picture in a more self-consistent manner.
Acknowledgments This work has been supported by NASA under grant NAG5-
12878 and contract NAS5-38099. I thank W. Abbett for a critical reading. I am
also grateful to Rob Rutten for LaTeX instruction, and to Bart de Pontieu for
meticulous keyboard entry.
Appendix: plasma parameters
The lower solar atmosphere marks the transition layer between regions of strik-
ing physical differences, and as one goes further up in height the tools of plasma
physics should become more important. This Appendix evaluates for conve-
nience several basic plasma-physics parameters for the conditions of the staple
VAL-C atmospheric model (Vernazza et al. 1981)1. This is a “semi-empirical
model” in which interprets a set of observations in terms of the theory of ra-
diative transfer, but without any effort to have self-consistent physics. Such a
model can accurately represent the spectrum but may or may not provide a
good starting point for physical analysis. Because the optical depth of a spec-
tral feature is the key parameter determining its structure, one often sees the
model parameters plotted against continuum optical depth τ5000 evaluated at
5000Å. Just for illustration, Figure 11 shows the VAL-C temperature separately
as a function of height, column mass, and optical depth. Note that features
prominent in one display may appear to be negligible in another
The VAL-C model is an “average quiet Sun” model, and like all static 1D
models, it cannot describe the variability of the physical parameters that theory
The VAL-C parameters are available within SolarSoft as the procedure VAL C MODEL.PRO.
18 Hudson
10000 1000 100 10
Height, km
10-10 10-8 10-6 10-4 10-2 100
Continuum opacity
10-8 10-6 10-4 10-2 100 102
Column mass, g/cm3
Figure 11. Illustration of the structure of a semi-empirical model, using
three different independent variables: the VAL-C temperature plotted against
height, optical depth, and column mass.
and observation require (see other papers in these proceedings, e.g., Carlsson’s
review). Thus we should regard the plasma parameters estimated here as order-
of-magnitude estimates only and note especially that the vertical scales, which
depend in the model on the inferred optical-depth scale, may be systematically
displaced.
The VAL-C model explicitly does not represent a chromosphere perturbed
by a flare. Vernazza et al. (1981) and many other authors give more appropriate
models derived by similar techniques for flares as well as other structures. As
the discussion of the γ-ray signatures in Section 10 suggests, though, a pow-
erful flare may be able distort the lower solar atmosphere essentially beyond
all recognition (especially in the impulsive phase). To estimate representative
plasma parameters I have therefore chosen just to start with the basic VAL-C
model, and we simply assume constant values of B at 10 G and 1000 G. The
actual magnetic field may vary through this region (the “canopy”) but the de-
tails are little-understood. The γ-ray literature usually uses a parametrization
of the magnetic field strength B ∝ Pαg (Zweibel & Haber 1983), where Pg is the
gas pressure.
The most complicated behavior of the plasma parameters happens prefer-
entially near the top of the VAL-C model range (for example, Figure 12 shows
that the collision frequency decreases below the plasma and Larmor frequencies)
above the helium ionization level (or even below this level for strong magnetic
fields). Because VAL-C ignores time dependences and 3D structure, and as-
sumes Te = Ti, we can expect that it has diminished fidelity as one approaches
the unstable transition region; thus one should be especially careful not to take
these approximations too literally. The following notes correspond to each panel
of the figure. Most of the plasma-physics formulae used in this Appendix are
from Chen (1984).
Chromospheric Flares 19
Figure 12. Various plasma parameters in the VAL-C model. We have as-
sumed representative B values of 10 G and 1000 G. The different panels show
the following, left to right and top to bottom: (a) Temperature. (b) Den-
sities: solid, total hydrogen density; dotted, electron density; dashed, He I
density; dash-dot, He II density. (c) Plasma beta: solid, for 1000 G; dotted,
for 10 G; light solid, electron density as a fraction of total hydrogen den-
sity; dash-dot, the plasma parameter. (d) Frequencies. Solid, electron and
ion plasma frequencies; dotted, electron gyrofrequencies for 10 and 1000 G;
dashed, electron and ion collision frequencies; dash-dot, electron/neutral col-
lision frequency. (e) Velocities: Solid, electron and ion thermal velocities;
dashed, Alfvén speeds for 10 and 1000 G. (f) Scale lengths: solid, electron
Larmor radii for 10 G and 100 G; dotted, Debye length; dashed, ion and
electron inertial lengths.
Temperature: The VAL-C model, like all of the semi-empirical models, sets
Te = Ti. It therefore cannot support plasma processes dependent upon dif-
ferent ion and electron temperatures, or more complicated particle distribution
functions (e.g. Scudder 1994).
Densities: Total hydrogen density, electron density, and densities of He I and
He II.
Dimensionless parameters: We approximate the plasma beta as
2(nH + 2ne)kT
B2/8π
with nH the hydrogen density, ne the electron density, Figure 12(c) gives the
number of electrons in a Debye sphere as the “plasma parameter” Λ.
20 Hudson
Frequencies: The plasma frequency, the electron and proton Larmor frequen-
cies, and the electron and ion and collision frequencies
νei = 2.4 × 10
−6nelnΛ/T
eV ; νii = 0.05 × 4νei; νeH = (nH/ne)νe
with ne in cm
−3, TeV the temperature in eV, using Z = 1.2 and the Coulomb
logarithm lnΛ = 23 - ln(n0.5e T
eV ) (Chen 1984; De Pontieu et al. 2001).
Note that the collision frequencies are small compared with the plasma and
Larmor frequencies above about 1000 km in this model. This means generally
that plasma processes must have strong effects on the physical parameters of
the atmosphere in this region.
Velocities: Electron and proton thermal velocities; Alfvén speeds vA for 10 and
1000 G.
Scale lengths: Electron Larmor radii for 10 and 1000 G, the ion inertial length
c/ωpi, the electron inertial length c/ωpe, and the Debye length λD. The iner-
tial lengths determines the scale for the particle demagnetization necessary for
magnetic reconnection. For VAL-C parameters the ion inertial length increases
to about 100 m in the transition region.
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