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Department of physics
Seminar - 4th class
Solar flare
Author: Neža Sušnik
Mentor: Prof. Tomaž Zwitter, Ph.D.
Ljubljana, March 2011
Abstract
In this seminar one of the most frequent phenomenon on the Sun, Solar flare is
described and explained. The structure of the Sun, its features and its magnetic field are also
dealt with. Solar flares are sometimes dangerous to the Earth, but their effect on the Earth's
ionosphere is spectacular to us. In last year the Sun's activity increased and so has our
interest in the Sun.
Contents
1 Introduction
1
2 The Sun
2
2.1 Solar interior
2
2.2 Solar atmosphere
3
2.3 Solar magnetism
6
3 Solar flare
7
3.1 Solar flare classification
9
3.2 The standard model of Solar flare
10
3.3 Magnetic reconnection
11
3.4 Observations
12
4 Conclusion
13
5 References
14
1 Introduction
The first Solar flare to be observed was on September 1, 1859 by the British
astronomer Richard Carrington. The event is the Solar storm of 1859; it is also named the
“Carrington event”. The result of the flare was a coronal mass ejection that sent charged
particles streaming toward the Earth, reaching the atmosphere only 18 hours after the
ejection. Once the particles reached the Earth they caused auroras down to tropical latitudes
such as Cuba or Hawaii, and set telegraph systems on fire. [1]
A Solar flare is an explosion on the Sun that happens when energy stored in twisted
magnetic fields, usually above sunspots, is suddenly released. Flares produce a burst of
radiation across the electromagnetic spectrum, from radio waves to x-rays and gamma-rays,
including high-energy particles, i.e. cosmic rays.[2] The intense radiation from a solar flare
travels to the Earth in eight minutes and can disrupt long distance radio signals or disturb a
satellite's orbit around the Earth. Energetic particles accelerated in Solar flares that escape
into interplanetary space are dangerous to astronauts and to electronic instruments in the
space.
1
The transient phenomena occurring in
the Solar atmosphere can be grouped together
under the term Solar activity: sunspots and
faculae occur in the photosphere; flares and
spicules belong to the chromosphere; and
prominences and coronal structures develop in
the corona. All Solar activity phenomena are
connected in this way or another with the 11
and 22-year sunspot cycle. To explain such a
phenomenon as a Solar flare one must first
introduce the sheer nature of the Sun. [3]
2 The Sun
The Sun is the largest object in our
Figure 1: A minor flare and a coronal mass ejection
Solar system and contains approximately
producing a prominence that rose up and out in a
99.8% of the total Solar system mass. It has a
curving arch on September 8th, 2010.[16]
diameter of about 1,390,000 km and its mass
is 1.9891x1030 kg. The Solar mass is strongly concentrated towards the centre. The outer
convection zone, which extends from immediately below the surface down to r/r ʘ ≈ 0.7,
contains only about 1.67% of the total mass. And only about 10% of the mass lies outside
r/rʘ ≈ 0.5, although 7/8 of the volume is there. At present, the Solar matter is approximately
73.5% hydrogen, 24.8% helium and 1.7% metals by mass fraction however chemical
composition changes with time due to thermonuclear reactions in the Solar core.[4]
The Sun does not have a definite boundary and in its outer parts the density of its
gases drops exponentially with increasing distance from its centre. Nevertheless, it has a welldefined interior structure. The Sun's radius is measured from its centre to the edge of the
photosphere. This is simply the layer above which the gases are too cool or too thin to
radiate a significant amount of light, and therefore the surface is most readily visible to the
naked eye.
2.1 The Solar interior
The Solar interior is not directly observable, and the Sun itself is opaque to visible light.
However, just as seismology uses waves generated by earthquakes to reveal the interior
structure of the Earth, the discipline of helioseismology makes use of pressure waves,
infrasound, traversing the Sun's interior to measure and visualize the star's inner structure.[6]
The core of the Sun is considered to extend from the center to about 20-25% of the
Solar radius. The temperatures and densities range between 15x10 6 K and 160 g/cm3 at the
center to 7x106 K and 20 g/cm3 at the outer edge of the core. It is the only region in the Sun
2
that produces an appreciable amount of thermal energy through fusion. The rest of the star is
heated by energy that is transferred outward from the core. [5]
From about 0.25 to about 0.7
Solar radii is the radiative zone. It is
a region of highly ionized gas and
the energy transport is primarily by
photon diffusion. The radiation
takes about 1 million years to find
its way out of the radiative layer,
even travelling at the speed of light,
due to collisions between the light
and the matter within the radiative
layer. In this zone the material gets
cooler from 7 to about 2 million
kelvin with increasing altitude, this
temperature gradient is less than
the value of the adiabatic lapse rate
and hence cannot drive convection.
Figure 2: Schematic view of the structure of the Sun and modes of
[5]
outward flow of energy.[5]
In the Sun's outer layer, from
its surface down to approximately 200,000 km (70% of the Solar
radii), is the convective zone. The high opacity makes it difficult for
photon radiation to continue outward and steep temperature
gradients are established which lead to convective currents. The
temperature drops from 2 million kelvin to about 6000 kelvin on the
outer layer. [6]
2.1 Solar atmosphere
Observationally, the outer Solar layers following the
convective zone have been divided into four spherically symmetric
layers
the
photosphere (not part
of the atmosphere),
chromosphere,
the
Figure 3: Sun's inner
transition region and temperature and density
profile. [19]
corona
lying
successively above one another. The last three
layers are part of the Solar atmosphere.
Figure 4: Sunspots and Solar granules on the
photosphere.[3]
The photosphere which is a thin layer only
several 100 km thick represents the surface of
the Sun. As the bubbles of upwelling, hot
3
plasma from convective zone, reach the surface of the photosphere, bright spots or granules
are created. The brighter spots are Solar granules and the large dark spots are sunspots.
Each granule measures about 1,000 km across, is as deep as the photosphere, and has a
lifetime of between 5 and 10 minutes. Each granule forms the topmost part of a Solar
convection cell. Spectroscopic observation of the photosphere within and around the bright
regions shows direct evidence for the upward motion of gas as it "boils" up from within. This
evidence proves that convection really does occur at or below the photosphere. Spectral lines
detected from the bright granules appear slightly bluer than normal, indicating Dopplershifted matter coming toward us with a velocity of about 1 km/s. Conversely, spectroscopes
focused on the darker portions of the granulated photosphere show the same spectral lines to
be redshifted, indicating matter moving away. The brightness variations of the granules result
strictly from differences in the temperature. The upwelling gas is hotter and therefore emits
more radiation than the cooler downwelling gas. The adjacent bright and dark gases appear
to contrast considerably, but in reality their temperature difference is less than about 500 K.
[5], [6]
Careful measurements also reveal a much larger-scale flow on the Solar surface.
Supergranulation is a flow pattern quite similar to granulation except that supergranulation
cells measure some 30,000 km across. As with granulation, material upwells at the centre of
the cells, flows across the surface, and then sinks down again at the edges. Scientists believe
that supergranules are the imprint on the photosphere of a deeper tier of large convective
cells.
We are most familiar with the photosphere because it is the visible surface of the Sun
and it produces most of the white light we see. According to the Stefan-Boltzmann law each
square meter of the Solar surface having the temperature T emits, in all directions, the light
of σT4 joules per second. Subsequently, the total emission of the Sun in one second, i.e. the
luminosity, equals
L ʘ = 4  R2ʘ  T 4 ≃ 3.845×10 26 W.
(1)
This fundamental relation also determines the effective temperature of the Sun when its
luminosity and radius are known. [3]
For a few seconds, just after the beginning and before the end of a total eclipse, the
Solar limb presents a most colourful view, the chromosphere
(“coloured sphere”). The chromosphere is a thin layer of the
Sun's atmosphere just above the photosphere, roughly 2,000 km
high. Although thin, its density decreases over almost seven
orders of magnitude from a high of 2x10 -7 g/cm3 at its boundary
with the photosphere, and decreasing to a low of 1x10 -14 g/cm3
where it merges with the Solar transition region. The
spectrograph reveals the chromosphere's “flash spectrum” which
shows a large number of lines that are dark in the normal Solar
spectrum. A prominent example is the red Hα line at 656.3 nm.
[4]
Figure 5: Total eclipse of 1999.
[6]
4
The large optical thickness of the chromospheric lines allows us to observe the
chromosphere on the disc, by means of a narrow band filter. Even more than the limb
observations, disc filtergrams demonstrate the extreme inhomogeneity of the Solar
chromosphere. The whole Sun is covered by numerous dark mottles. The dark mottles more
or less outline a network that follows the supergranulation pattern. Since the magnetic flux is
continuously convected towards the supergranular boundaries, it is well possible that the
mottles owe their existence to local enhancements of the Solar magnetic field. These mottles
must be closely related to the spicules seen at the limb, although a unique identification with
spicules is difficult. [4]
Spicules are long, thin spikes of
matter that leave the Sun's surface at typical
velocities of about 100 km/s, reaching
several thousand kilometres above the
photosphere. Spicules are not spread evenly
across the Solar surface, instead, they cover
only about 1% of the total area, tending to
accumulate
around
the
edges
of
supergranules. [4]
There is even more to the
chromospheric phenomenon: the outwards
increasing temperature. Considerably hotter
than the photosphere, the chromosphere is
heated by hydromagnetic waves and
compression waves originated by spicules
and granules. The temperature rises from
about 6,000 K to about 20,000 K.[3]
The layer between the relatively cool
chromosphere and the hot corona is called
the transition region. This is more like a
temperature regime rather than a geometric
layer not only because of the extreme
spatial inhomogeneity of this region, but
also because the transition is so sharp that Figure 6: Skylab measured the temperature (solid curve)
there is virtually a discontinuity in the
and density (dashed curve) of the chromosphere
between
the thinner transition region and the lower
temperature. Within this region the
photosphere
(darker orange). [20]
temperature rapidly increases from 20,000 K
to 1,000,000 K. This phenomenon is called the temperature catastrophe and is a phase
transition analogous to boiling water to make steam. [5]
The corona is the Sun's outer atmosphere. It is visible during total eclipses of the Sun
as a pearly white crown surrounding the Sun. The corona displays a variety of features
including streamers, plumes, and loops. These features change from eclipse to eclipse and
the overall shape of the corona changes with the sunspot cycle. The coronal gases are super5
heated to temperatures higher than 1,000,000 K. At these high temperatures both hydrogen
and helium are completely stripped of their electrons. Even minor elements like carbon,
nitrogen, and oxygen are stripped down to bare nuclei. Only the heavier trace elements like
iron and calcium are able to retain a few of their electrons in this intense heat. It is the
emission from these highly ionized elements that produces the spectral emission lines that
were so mysterious to early astronomers. [6]
The corona shines brightly in x-rays
because of its high temperature. On the
other hand, the "cool" Solar photosphere
emits very few x-rays. This allows us to view
the corona across the disk of the Sun when
the Sun in X-rays is observed. The corona is
not always evenly distributed across the
surface of the Sun. During periods of quiet,
the corona is more or less confined to the
equatorial regions, with coronal holes
covering the polar regions. However during
the Sun's active periods, the corona is evenly
distributed over the equatorial and polar Figure 7: The x-ray corona showing coronal holes and
regions, though it is most prominent in the
coronal bright spots.[5]
areas with the sunspot activity. [6]
2.2 Solar magnetism
The Sun's magnetic field is (widely believed to be) generated by a magnetic dynamo
within the Sun and it changes dramatically over the course of just a few years, and the fact
that it changes in a cyclical manner indicates that the magnetic field continues to be
generated within the Sun. The electric currents that are generated within the Sun,
presumably in the transition layer, by the flow of the Sun's hot, ionized gases, produce
magnetic fields.[7]
Magnetic fields within the Sun are stretched out and wound around the Sun by
differential rotation. This is called the omega-effect after the Greek letter used to represent
rotation. The Sun's differential rotation with latitude can take a north-south oriented magnetic
field line and wrap it once around the Sun in about 8 months. [7]
Twisting of the magnetic field lines is caused by the effects of the Sun's rotation. This
is called the alpha-effect. Recent dynamo models assume that the twisting is due to the effect
of the Sun's rotation on the rising "tubes" of magnetic field from deep within the Sun. The
twist produced by the alpha effect makes sunspot groups that obey Joy's law (tilt of the
sunspot groups) and also makes the magnetic field reverse from one sunspot cycle to the
next one (Hale' law). [7]
Sunspots are the most prominent magnetic feature on the Solar surface (sometimes
they can be seen with the naked eye). The number of sunspots visible on the Solar surface
changes fairly regularly in time. The number of sunspots increases and decreases over a
6
period of 11 years. This is called the Solar cycle. Moreover, the position of sunspots also
changes in time. In 11 years sunspots move from high latitudes to the equator. By plotting
the position of the sunspots versus time the so called butterfly diagrams is obtained. [8]
Figure 8: The butterfly diagram presents the latitude of sunspots with time. This one also indicates the magnetic
field reverse from one sunspot cycle to the next.[8]
Sunspots appear as dark spots on the surface of the Sun. Temperatures in the dark
centres of sunspots drop to about 3700 K (compared to 6000 K for the surrounding
photosphere). They occur where magnetic fields suppress convection of hot matter to the
surface. They typically last for several days, although the very large ones may live for several
weeks. Sunspots are magnetic regions on the Sun with magnetic field strengths of up to 0.20.3 T. The field is strongest in the darker parts of the sunspots - the umbra, and the field is
weaker and more horizontal in the lighter part - the penumbra. Sunspots usually come in
groups with two sets of spots. One set will have positive or north magnetic field while the
other set will have negative or south magnetic field. These sets of north-south magnetic field
sunspots are usually oriented east to west, and the orientation on northern hemisphere is
always opposite to the orientation on the southern hemisphere.[8]
3 Solar flare
The total average energy output of the outer layers of the Sun is ≈ 10-4 of the
photospheric radiation. The local and short-lived enhancements may exceed that average by
a factor 103 to 104, for the largest events the flux of energy may exceed the photospheric
flux. The explosive energy release is in a totally different form. A variety of observational
techniques has therefore been employed, and the phenomena seen have been given various
names. Sometimes these phenomena are observed in isolation, but more often, in particular
for the largest events, all or many of them are associated with each other. Flares occur near
7
sunspots, usually along the dividing line (neutral line) separating opposite magnetic fields.
[4]
A Solar flare may be naively defined
as a rapid brightening in Hα, but it can
simultaneously have manifestations right
across the electromagnetic spectrum and
may eject high energy particles and blobs
of plasma from chromosphere out into the
Solar wind. The Hα flare is formed in the
chromosphere and has two basic stages.
During the flash phase, which lasts
typically 3 minutes, but sometimes an
hour, the intensity and area of the
emission rapidly increase in value. Then,
in the main phase, the intensity slowly
declines over about an hour or
occasionally as long as a day. [9]
Above the cool Hα flare there lies a
high-temperature coronal region, which
may be heated during the flare to tens of
millions of degrees. It exhibits variations
co-incident with the flash and main phase Figure 9: A schematic profile of intensity for a typical flare
in several wavelengths.[9]
but also shows two more distinct phases.
As seen in figure 7, the soft X-ray emission (<10 keV) possesses a pre-flare phase for minutes
before flare onset, because of an enhanced thermal emission from the coronal plasma. Also,
for 10-100 seconds at the start of the flare, an impulsive phase is sometimes, as indicated by
the appearance of the microwave burst and a hard X-ray burst (>30 keV), caused by highly
accelerated electrons. These highly accelerated electrons are accelerated nearby the magnetic
reconnection point and travel down along the two legs of a coronal loop into the low corona
and chromosphere, where they heat the plasma very rapidly (in figure 11). The hard X-ray
exhibits a complex spiky structure, with the smallest real time-scales of about 2 seconds in
moderate events and 10 seconds in the large ones. It is these which give the time-scale for
electron acceleration during the impulsive phase. After the impulsive phase some of the
largest events show a distinct second hard X-ray component due to a second phase of
particle acceleration. For other events, the impulsive phase may take place after the Hα
intensity has started increasing or it may be absent. In the latter case, where there is little
particle acceleration, the events are known as thermal flares, they tend to occur in less
complex regions and have a slow rise to flare maximum.[9]
8
3.1 Solar flare classification
Flares are classified in a multitude of different ways, but there are two main types of
flares, which appear to require quite different physical mechanisms.
Simple-loop flare or compact flare:
Most flares and subflares are of this type. It is a small flare, in which essentially a
single magnetic loop or flux tube brightens in X-rays and remains apparently unchanged in
shape and position throughout the event. The loop may have a structure consisting of several
loops and may cause a simple brightening in Hα at the feet of the loop. The flare can occur
within a large-scale unipolar region or near a simple sunspot, where little excess magnetic
energy is stored. Occasionally, it may be accompanied by a surge as a stream of plasma, with
an average density of 1016 particles/m3, which is squirted upwards.[9]
Two-ribbon flare:
All major events are of this type. It is much larger than a compact flare and takes
place near a Solar prominence, a loop of plasma confined between two magnetic field lines,
which shows up in Hα pictures as a dark ribbon called filament. When the filament is located
in part of the quiet Sun, such as a remnant active region, the flare tends to be slow, longlived and not very energetic, presumably because the magnetic field is relatively week near
such a quiescent filament. Active-region filaments are located near the intense and
sometimes complex field of sunspots and they are associated with the most violent and
energetic flares. During the flash phase of two-ribbon flare, two ribbons of Hα emission form,
one on each side of the filament and throughout the main phase, the ribbons move apart at 2
to 10 km/s. Frequently, they are seen to be connected by an arcade of “post-flare” loops.
Occasionally, the filament remains intact, though slightly disturbed, but usually it rises and
disappears completely. Such an eruption of the filament begins slowly in the pre-flare phase,
typically 10 minutes, but up to an hour, before the flare onset, and continues at the flash
phase with a much more rapid acceleration than before.[9]
There are many differences between simple-loop
and two-ribbon events. Simple-loop flares tend to have at
most a single hard X-ray spike lasting about a minute,
whereas two-ribbon flares may have multiple spikes. In
soft X-rays simple-loop flares are characterised by small
volumes, low heights, large energy densities and short
time-scales, while two-ribbon flares have the opposite
properties and often produce coronal transients. Also the
energy release may be confined to the impulsive phase of
a simple-loop event, but it continues throughout the main
phase of a two-ribbon event.[9]
Figure 10: A typical example of a tworibbon flare.[21]
Other classifications of Solar flares are in two ways: Hα classification and Soft X-ray
classification. The latter classifies flares according to the order of the magnitude of the peak
burst intensity measured ate the Earth in the 0.1 to 0.8nm wavelength band as follows:
9
Hα classification
Class
Area
(Sq.
Deg.)
Soft X-ray class
Radio flux
Peak flux
at 5000
Class
in 1 - 8 Å
MHz
2
(W/m )
S
2
5
A
10-8-10-7
1
2.0-5.1
30
B
10-7-10-6
2
5.2-12.4
300
C
10-6-10-5
3
12.5-24.7
3000
M
10-5-10-4
4
>24.7
3000
X
>10-4
A multiplier is used to indicate the level within each class: M6 = 6 x 10 -5 W/m2. [10]
3.2 The standard model of a Solar flare
The standard model of a Solar flare consists of eight steps:
1. Magnetic free energy is stored in the corona, due to either motions of the
photospheric footpoints of loops or to the emergence of current-carrying field from below the
photosphere.
2. A cool, dense filament forms, suspended by the magnetic field, over the neutral line.
3. The field evolves slowly through equilibrium states, finally reaching non-equilibrium
which causes the closed field to rise and erupts outward.
4. The reconnection of the field
below the rising filament provides
plasma heating and particle acceleration
that is called the flare. The region at the
reconnection is also the hottest area of
the solar flare.
5. The accelerated particles follow
the field lines and interact with the
chromosphere, heating it and causing
the evaporation up-flows of the plasma.
6. “Post-flare loops” are formed
over the neutral line; they gradually cool
down by radiation.
7. More field lines are involved in
reconnection; the reconnection site is
going up, forming new post-flare loops
situated above the previously created
ones.
8. A multi-temperature arcade is
formed with “older” cooler loops being
Figure 11: Schematic view of Solar flare process. [10]
below “new” hotter loops.[10]
10
3.3 Magnetic reconnection
It is generally accepted that the
energy released during Solar eruptions
(flares,
coronal
mass
ejections,
prominence eruptions) is stored in the
magnetic field before the eruption.
Theoretical models of Solar eruptions
invariably
include
magnetic
reconnection as a physical process for
the release of magnetic energy and its
conversion into other forms of energy
such as bulk flow energy, thermal Figure 12: The basic configuration of two-dimensional steadyenergy, and non-thermal kinetic
state reconnection.[12]
energy.
Magnetic reconnection is a fundamental dynamical process in highly conductive
plasmas. The basic equations for a steady state MHD specialized for a resistive
magnetohydrodynamics, using the incompressibility condition ( ∇·v = 0) with constant density
ρ0 are:
(2)
(3)
(4)
(5)
(6)
 v⋅∇ v = −∇ p  j×B
E 0  v× B⋅e z = j z /
∇⋅v = 0
 ∇× B⋅e z = 0 j z
∇⋅B = 0
There is no fully satisfactory analytical treatment of the system of equations (2)–(6). There
are solutions for the external (ideal) region and solutions for the diffusion region based on
singular asymptotic expansions. From a set of control parameters ( ρ0, p0, v0, B0, L (the global
length scale) and σ (index 0 denotes values at (x0,0) from figure 11)) to disregard
configurations that are merely the result of a similarity transformation, there are three
independent dimensionless quantities formed, which are conveniently chosen as
M0 =
v0
,
a0
S0 =
a0 L
,

0 =
2  0 p0
2
B0
(7)
where a0 is the (inflow) Alfvén velocity B0/√μ0ρ0. The parameter β0 is ignored if β0 is negligibly
small or if the pressure is constant in the external region. (It is only the gradient of the
pressure that counts.) Then, the reconnection is a two-parameter process described by M0
and S0. The parameter M0 is usually called the reconnection rate. It measures the velocity
with which the plasma enters the region of consideration (normalized by the local Alfvén
velocity).[12]
In the Sweet-Parker (P. A. Sweet and E. N. Parker) model it is assumed that the
diffusion region is a thin extended structure. The external region is largely homogeneous such
11
that approximately B and S are constant outside the diffusion region. Under these conditions,
the reconnection rate is too low to be relevant for typical conditions in stellar atmospheres
and space plasmas. This model gives magnetic reconnection time of several days.[12]
In Petschek's (H. E. Petschek) model it is assumed that the diffusion region is very
small in length and therefore even smaller in width. In this case, it is necessary to consider
the presence of slow-mode shock waves that accelerate the plasma to Alfvén speed and
convert most of the magnetic energy into kinetic energy and heat. Typically this reconnection
rate is considerably larger than that of the Sweet–Parker process.[12]
Several authors have generalized the models by Sweet and Parker and by Petschek in
various respects. The most general are the fast reconnection models of Priest and Forbes.
They included electrical currents in the external region and obtained a description that
contained the Sweet-Parker and Petschek models as particular cases.[12]
The original meaning of reconnection is a breakdown of magnetic field line
conservation. In this process strong electric fields are created that accelerate charged
particles. The notion of general magnetic reconnection to occur is if
∫ E⋅ds
≠ 0
(8)
where the integral is evaluated for field lines passing through a localized non-ideal region
embedded in an otherwise ideal plasma. The criterion is sufficient for a breakdown of
magnetic line conservation, provided all magnetic field lines start and end in the ideal region
outside non-ideal region. This is a consequence of the general form of magnetic field line
conservation
∂B
− ∇ × w× B =  B
∂t
(9)
where w is the transport velocity of the field lines, which can be identified with the plasma
velocity v in the ideal region but may differ from it in non-ideal processes.[12]
3.4 Observations
The following spacecraft missions have observed the Sun and its atmosphere for flares
and other activities:
Yohkoh was a Solar observatory spacecraft of the Institute of Space and Astronautical
Science (Japan) with United States and United Kingdom collaboration. It observed the solar
atmosphere in X-ray radiation continuously for more than ten years. The mission ended after
more than ten years of successful observations in 2001 and after four years of inactive state
burned up during re-entry over South Asia in 2005. [13]
SOHO, the Solar and Heliospheric Observatory, is a project of international
collaboration between ESA and NASA to study the Sun from its deep core to the outer corona
and the Solar wind. SOHO has studied the Sun-Earth interaction since 1995.[14]
12
RHESSI or Reuven Ramaty High Energy Solar
Spectroscopic Imager is the sixth mission in the line of
NASA Small Explorer missions. Launched on 5 February
2002, its primary mission is to explore the basic physics of
particle acceleration and explosive energy release in Solar
flares. The primary scientific objective of RHESSI is to
understand the following processes that take place in the
magnetized plasmas of the Solar atmosphere during a
flare: impulsive energy release, particle acceleration, and
particle and energy transport. RHESSI was the first satellite
to image Solar gamma rays from a Solar flare and was the
first to accurately measure terrestrial gamma ray flashes Figure 13: A very large coronal mass
that come from thunder storms.[15]
ejection on December 2, 2002, imaged
HINODE is a Japanese project in partnership with
by SOHO.[13]
NASA to study the Sun, explore the magnetic fields of the
Sun, and expand our understanding of the mechanisms that power the Solar atmosphere and
drive Solar eruptions. [16]
SDO or Solar Dynamics Observatory is the first mission to be launched for NASA's
Living With a Star Programme, a programme designed to understand the causes of Solar
variability and its impacts on the Earth. SDO's goal is to understand the Solar variations that
influence life on the Earth and humanity's technological systems by determining how the
Sun's magnetic field is generated and structured, and how this stored magnetic energy is
converted and released into the heliosphere and geospace in the form of Solar wind,
energetic particles, and variations in the Solar irradiance. SDO was lunched on February 11,
2010.[17]
STEREO, Solar TErrestrial RElations Observatory, provides a unique views of the SunEarth system. The satellites trace the flow of energy and matter from the Sun to Earth as well
as reveal the 3-D structure of coronal mass ejections and help us understand why they
happen. STEREO also provides alerts for Earth-directed Solar ejections, from its unique sideviewing perspective adding it to the fleet of Space Weather detection satellites. [18]
4 Conclusion
At the beginning of this year the 24 th Solar cycle was visibly entered. The Sun's activity
has increased and there are many Solar flares every week. For example, on February 15, an
X2 flare erupted and was seen by SDO in extreme ultraviolet light. This was the largest Solar
flare since December 5, 2006. Because of the increased Sun's activity and because of the
frightening (elusive) predictions that there will be a massive Solar storm in a near future,
Sun's activities are getting more attention every day. The Sun's activity has also an effect on
the Earth's weather, because a prolonged period of change of the energy reaching the Earth
would result in a change in climate, but short term changes in the energy from the Sun are
rarely perceptible.[18]
13
Although Solar flares have been discovered more than 150 years ago, the exact
mechanism of Solar flares is still not known. To understand Solar flares, plasma and Solar
magnetism have to be understood better. This is the reason why all satellites observing the
Sun are constantly sending new data to the Earth and new information is tried to be
processed.
5 References
[1] http://en.wikipedia.org/wiki/Solar_storm_of_1859 (3.3.2011)
[2] http://en.wikipedia.org/wiki/Solar_flare (3.3.2011)
[3] http://neutrino.aquaphoenix.com/un-esa/sun/sun-chapter1.html (3.3.2011)
[4] M. Stix, The Sun, An introduction (Springer-Verlag, Berlin, 1989)
[5] http://stargazers.gsfc.nasa.gov/resources/amazing_structure.htm (4.3.2011)
[6] http://en.wikipedia.org/wiki/Sun (3.3.2011)
[7] http://solarscience.msfc.nasa.gov/dynamo.shtml (4.3.2011)
[8] http://solarscience.msfc.nasa.gov/SunspotCycle.shtml (4.3.2011)
[9] E. R. Priest, Solar flare magnetohydrodynamics (Gordon and Breach science
publisher, 1981)
[10]
http://www2.warwick.ac.uk/fac/sci/physics/teach/module_home/px420/handouts/magr
ec_new.pdf (7.3.2011)
[11] https://www.cfa.harvard.edu/~scranmer/ITC/eaaa_reconn_schindler.pdf
(7.3.2011)
[12] http://en.wikipedia.org/wiki/Yohkoh (8.3.2011)
[13] http://sohowww.nascom.nasa.gov/home.html (8.3.2011)
[14] http://hesperia.gsfc.nasa.gov/hessi/sheet.htm (8.3.2011)
[15] http://www.nasa.gov/mission_pages/hinode/index.html (8.3.2011)
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[16] http://www.nasa.gov/mission_pages/sdo/main/index.html (8.3.2011)
[17] http://stereo.gsfc.nasa.gov/mission/concept.shtml (8.3.2011)
[18] http://www.newton.dep.anl.gov/askasci/wea00/wea00141.htm (21.3.2011)
[19] http://www.fas.org/irp/imint/docs/rst/Sect20/A5a.html (21.3.2011)
[20] http://en.wikipedia.org/wiki/Photosphere (3.3.2011)
[21] http://www.kwasan.kyoto-u.ac.jp/general/facilities/dst/index_en.html (8.3.2011)
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