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 Solar system and contains approximately 99.8% of the total Solar system mass. It has a 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/rʘ ≈ 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 well-defined 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 15x106 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

Figure 1: A minor flare and a coronal mass ejection producing a prominence that rose up and out in a

curving arch on September 8th, 2010.[16]

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.[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 transition region and corona - lying successively above one another. The last three layers are part of the Solar atmosphere.

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

Figure 2: Schematic view of the structure of the Sun and modes of outward flow of energy.[5]

Figure 3: Sun's inner temperature and density

profile. [19]

Figure 4: Sunspots and Solar granules on the photosphere.[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 Doppler-shifted 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

Lʘ = 4 Rʘ2 T 4 ≃ 3.845×1026W. (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]

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Figure 5: Total eclipse of 1999. [6]

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 there is virtually a discontinuity in the temperature. Within this region the 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 super-

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Figure 6: Skylab measured the temperature (solid curve) and density (dashed curve) of the chromosphere

between the thinner transition region and the lower photosphere (darker orange). [20]

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 regions, though it is most prominent in the 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

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Figure 7: The x-ray corona showing coronal holes and coronal bright spots.[5]

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]

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 0.2-0.3 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

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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, 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 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

Figure 9: A schematic profile of intensity for a typical flare in several wavelengths.[9]

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, long-lived 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]

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:

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Figure 10: A typical example of a two-ribbon flare.[21]

Hα classificationRadio fluxat 5000

MHz

Soft X-ray class

ClassArea(Sq.Deg.)

ClassPeak fluxin 1 - 8 Å (W/m2)

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 below “new” hotter loops.[10]

10

Figure 11: Schematic view of Solar flare process. [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 energy, and non-thermal kinetic 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:

v⋅∇ v = −∇ p j×BE0 v×B⋅e z = j z/

∇⋅v = 0∇×B⋅ez = 0 j z

∇⋅B = 0

(2)(3)(4)(5)(6)

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

M 0 =v0

a0, S0 =

a0 L

, 0 =20 p0

B02 (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

Figure 12: The basic configuration of two-dimensional steady-state reconnection.[12]

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∂ t

− ∇×w×B = B (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]

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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 that come from thunder storms.[15]

HINODE is a Japanese project in partnership with 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 Sun-Earth 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 side-viewing perspective adding it to the fleet of Space Weather detection satellites. [18]

4 Conclusion

At the beginning of this year the 24th 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

Figure 13: A very large coronal mass ejection on December 2, 2002, imaged

by SOHO.[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/magrec_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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