All around the world, everyone enjoy the day and night one after another. The day begins with Sunrise, and it ends with Sunset. Every day we are enjoying the natural play of the Sun. The Sun is giving light and warmness.



The size of the Suns’ disc on the sky covers 32 arc minutes, almost exactly the same as the angular size of the moon as soon from the Earth. Sun is a ball of hot gas in the plasma state with a mass of 1.9891x1030 kg, roughly 330,000 times the mass of the Earth and diameter of 1.392x106 km. (roughly 109 times the diameter of the Earth) [Woolfson M., 2000]. Since the Sun has volume roughly one million times that of the Earth. But the mass equivalent to only one third of a million of that of the Earth, about 1.4 times the density of water. But that average disguises a very wide range of densities, from a super dense core to the tenuous outer layer. The Sun is an ordinary star, roughly halfway through its life time. The Sun has a luminosity of 3.83x1026 watts. The Sun is located in the outskirts of the Milky Way, about 30,000 light years from its center. Although we would have to travel about 2,00,000 times the distance between the Sun and Earth, i.e., 2,00,000 AU (1 Astronomical Unit (AU) = distance of the Earth from the Sun).

The Sun’s stellar classification, based on spectral class, is G2V, and is informally designated as a yellow dwarf, because its visible radiation is most intense in the yellow-green portion of the spectrum and although its color is white, from the surface of the Earth it may appear yellow because of atmospheric scattering of blue light.

. In the spectral class label, G2 indicates its Surface temperature of approximately 5778 K (5505 ̊C), and V indicates that the Sun, like most stars, is a main-sequence star, and thus generates its energy by nuclear fusion of hydrogen nuclei into helium. In its core, the fuses 620 million metric tons of hydrogen each second.

Characteristics Features of Sun:

The sun is giant fusion reactor converting massive amount of hydrogen into helium. The sun produces vast amount of energy by thermonuclear reactions. The thermonuclear reactions takes place at very high temperature, where the hydrogen nuclei get converted into helium nuclei by the process of fusion reaction and releases vast amount of energy of an average of 4.25 million tons of hydrogen converted into helium each second in the sun is being consumed gradually and helium nuclei are being created. The sun’s structure can be described as a system of six main layers:

  • 1. Core,
  • 2. Radioactive zone,
  • 3. Convective zone,
  • 4. Photosphere,
  • 5. Chromospheres,
  • 6. Corona.


The core is the central region where thermonuclear fusion reactions, consume hydrogen to form helium [Garcia, et al. 2007]. These reactions release the energy as visible light and are highly sensitive to temperature and density. The matter in the core consists of plasma at the temperature of about 15,000,000°C and density of the core is about 150 g/cm3 [Basu, et al. 2009]. Both temperature and density decrease as one move outward from the centre of the Sun. The nuclear burning is almost completely shutoff beyond the outer edge of the core. In stars, like the sun, the nuclear burning (and energy release) takes place through a three step process called proton-proton or PP chain. In the first step two protons collide to produce deuterium, a positron and a neutrino. In the second step a proton collides with the deuterium to produce a helium-3 nucleus and the gamma ray. In the third step two helium-3 nuclei infuse to produce a normal helium-4 nucleus with the release of two protons and the energy equivalent of mass loss.

Radioactive Zone

The radioactive zone extended outwards, surrounds the core to the tachocline at the base of the convective zone is characterized by the method of energy transport radiation and helping it to maintain high temperature needed to sustain nuclear fusion. The energy generated in the core is carried by electromagnetic radiations (photons) that bounces from particle to particle through the radioactive zone although the photons take about a million years to finally reach the interface layer. The density varies from 20 g/cm3 (about density of gold) down to 0.2 g/cm3 (less than the density of water) from the bottom to the top of the radioactive zone. The temperature falls from 7000,000°C to about 2,000,000°C over the same distance.
The radioactive zone and the convection from a transition layer, the tachocline, a region where the sharp regime change between the uniform rotation of the radioactive zone and the differential rotation of the convection zone results in a large shear condition where successive horizontal layers slide past one another. The fluid motions found in the convection zone above disappear slowly from top to bottom of this layer, matching the calm characteristics of the radioactive zone on the bottom. Present hypothesis says that a magnetic dynamo within this layer generates the Sun’s magnetic field.

Convective Zone

In the Sun’s outer layer, from its surface down to approximately 200,000 km (or 70% of the solar radius), the solar plasma is not dense enough or hot enough to transfer the thermal energy of the interior outward through radiation; in other words it is opaque enough. As a result, thermal convection occurs as thermal columns carry hot material to the surface (photosphere) of the Sun. Once the material cools off at the surface, it plunges downward to the base of the convection zone, to receive more heat from the top of the radioactive zone. At the visible surface of the Sun, the temperature has dropped to 5,700 K and the density to only 0.2 g/m3.
The thermal columns in the convection zone form an imprint on the surface of the solar as the solar granulation and super-granulation. The turbulent convection of this outer part of the solar interior causes a “small-scale” dynamo that produces magnetic north and south poles all over the surface of the Sun.


The visible surface of the Sun, the photosphere, is the layer below which the Sun becomes opaque to visible light. About the photosphere visible sunlight is free to propagate into space, and its energy escapes the Sun entirely. The change in opacity is due to the decreasing amount of H- ions, which absorb visible light easily. Conversely, the visible light we see is produced as electrons react with hydrogen atoms to produce H- ions. The photosphere is tens to hundreds of kilometers thick, being slightly less opaque than air on Earth. Because the upper part of the photosphere is cooler than the lower part, an image of the Sun appears brighter in the center than on the edge or limb of the solar disk, in a phenomenon known as limb darkening [Abhyankar, 1977]. Sunlight has approximately a black-body spectrum that indicates its temperature is about 6,000 K, interspersed with atomic absorption lines from the tenuous layers above the photosphere

The photosphere has a particle density of ~1023 m-3 (this is about 0.37% of the particle number per volume of Earth’s atmosphere at sea level; however, photosphere particles are electrons and protons, so the average particle in air is times as heavy).


The coolest layer of the Sun is a temperature minimum region about 500 km above the photosphere, with a temperature of about 4100 K [Abhyankar, 1977]. This part of the Sun is cool enough to support simple molecules such as carbon monoxide and water, which can be detected by their absorption spectra [Solanki and Ayres, 1994]. Above the temperature minimum layer is a layer about 2000 km thick, dominated by a spectrum of emission and absorption lines. It is called the chromospheres from the Greek root chromo, meaning color, because the chromospheres is visible as a colored flash at the beginning and end of total eclipses of the Sun. The temperature in the chromospheres increases gradually with altitude, ranging up to around 20000K near the top. In the upper part of chromospheres helium becomes partially ionized [Hansteen and Leer, 1997].


Above the chromospheres, in a thin (~ 200 km) transition region, the temperature rises rapidly from around 20,000 K in the upper chromospheres to coronal temperature closer to 1x106K [Erdelyi, et al. 2007]. The temperature increase is facilitated by the full ionization of helium in the transition region, which significantly reduces radioactive cooling of the plasma. The transition region does not occur at a well-defined altitude. Rather, it forms a kind of nimbus around chromospheres features such as speckles and filaments, and is in constant, chaotic motion. The transition region is not visible easily from Earth’s surface, but is readily observable from space by instruments sensitive to the extreme ultraviolet portion of the spectrum [Hansteen and Leer, 1997].

Characteristics Features of Chromospheres

The Chromospheres Network

The chromospheres network is a web-like pattern most easily seen in the emission of the red line of hydrogen (H-alpha) and the ultraviolet line of calcium (Ca II K-from calcium atoms with one electron removed). The network outlines the super granule cells and is due to the presence of bundles of magnetic field lines.

Filament and Plage

Filaments are dark, thread like features seen in the red line of hydrogen (H-alpha). These are dense, somewhat cooler clouds of material that are suspended above the solar surface by loops of magnetic fields. Plage are bright patches surrounding sunspots that are also best seen in H-alpha. Plages are also associated with concentration of magnetic fields and form a part of the network of bright emissions that characterize the chromosphere.


Prominences are dense clouds of material suspended above the surface of the sun by loops of magnetic field. Prominences and filaments are actually the same things except that prominences are seen projecting out above the limb, or edge, of the sun. Both filaments and prominences can remain in a quiet state for days or a week. However, as the magnetic loops that support them slowly changes, filaments and prominences can erupt and rise-off of the sun over the course of a few minutes or hours.


Spicules are small, jet like eruptions seen throughout the chromospheric network. They appear as short dark streak in the H-alpha image. They last for a few minutes but in the process eject material-off of the solar surface and outward into the hot corona at high speeds.


The corona is the extended outer atmosphere of the Sun, which is much larger in volume than the sun itself. The corona continuously expands into space forming the solar wind, which fills all the Solar System. The low corona, near the surface of the Sun, has a particle density around 1015-1016m-3 [Hansteen and Leer, 1997]. The average temperature of the corona and solar wind is about 1-2x106K; however, in the hottest region it is 8-20x106K, while no complete theory yet exists to account for the temperature of the corona, at least some of its heat is known to be from magnetic reconnection [Erdelyi, et al. 2007].

The heliosphere, which is the cavity around the Sun filled with the solar wind plasma, extends from approximately 20 solar radii (0.1 AU) to the outer fringes of the Solar System. Its inner boundary is defined as the speed of Alfven waves. Turbulence and dynamic forces outside this boundary is defined as the layer in which the flow of the solar wind becomes super-alfvenic that is, where the flow becomes faster than the speed of Alfven waves [Emslie, et al. 2003]. Turbulence and dynamic forces outside this boundary cannot affect the shape of the solar corona within, because the information can only travel at the speed of Alfven waves. The solar wind travels outward continuously through the heliosphere, forming the solar magnetic field into a spiral shape, until it impacts the heliopause more than 50 AU from the sun. In December 2004, the Voyager 1 probe passed through a shock front that is thought to be part of the heliopause. Both of the Voyager probes have recorded higher levels of energetic particles as they approach the boundary.

Magnetic Field of Sun

The sun is a magnetically active star. It supports a, changing magnetic field that varies year-to-year and reverses direction about every years around solar maximum. The Sun’s magnetic field leads to many effects that are collectively called solar activity, including sunspots on the surface of the Sun, solar flares, and variations in solar wind that carry material through the Solar System. Effects of solar activity on Earth include auroras at moderate to high latitudes and the disruption of radio communications and electric power. Solar activity is thought to have played a large role and electric power. Solar activity is thought to have played a large role in the formation and evolution of the Solar System. Solar activity changes the structure of Earth’s outer atmosphere.


All matter in the Sun is in the form of gas and plasma because of its high temperatures. This makes it possible for the Sun to rotate faster at its equator (about 25 days) than it does at higher latitude (about 35 days near its poles). The differential rotation of the Sun’s latitudes causes its magnetic field lines to become twisted together over time, causing magnetic field loops to erupt from the Sun’s surface and trigger the formation of the Sun’s dramatic sunspots and solar prominences (see magnetic reconnection). This twisting action creates the solar dynamo and an 11 year solar cycle of magnetic activity as the Sun’s magnetic field reverses itself about every 11 year. The solar magnetic field extends well beyond the Sun itself. The magnetized solar wind plasma carries Sun’s magnetic field into the space forming what is the interplanetary magnetic field. Since the plasma can only move along the magnetic field lines, the interplanetary magnetic field is initially stretched radially away from the Sun. because the fields above and below the solar equator have different polarities pointing towards and away from the sun, there exists a thin current layer in the solar equatorial plane, which is called the heliospheric current sheet. At the large distances the rotation of the Sun twists the magnetic field and the current sheet into the Archimedean spiral like structure called the Parker spiral. The interplanetary magnetic field is much stronger than the dipole component of the solar magnetic field. The Sun’s 50-400 µT (in the photosphere) magnetic dipole field reduces with the cube of the distance to about 0.1 nT at the distance of the earth. However, according to spacecraft observations the interplanetary field at the Earth’s location is about 100 times greater at around 5 nT [Lodders, 2003].

Solar Activities



Sunspots are well-defined surface areas that appear darker than their surrounding because of lower temperatures. Sunspots are regions of intense magnetic activity where convection is inhibited by strong magnetic fields, reducing energy transport from the hot interior to the surface. The magnetic field causes strong heating in the corona, forming active regions that are the source of intense solar flares and coronal mass ejections. The largest sunspots can be tens of thousands of kilometers across. Why are sunspots dark? Basically the strong magnetic field, not allowing motions across the field lines, quenches convection inside the spot. Since convection is the main source of energy transport just below the surface, less energy reaches the surface through the spot. Sunspots looking darker than surrounding, more developed spots have dark interior, the umbra, surrounded, or at least partly so, by a lighter area known as the penumbra shown in Fig. 2.3.

The Sunspot Cycle or Solar Cycle


The number of sunspots visible on the Sun is not constant, but varies over an 11-year cycle known as the solar cycle. At a typical solar minimum, few sunspots are visible, and occasionally none at all can be seen. Those that do appear are at high solar latitudes. As the sunspot cycle progresses, the number of sunspots increases and they move closer to the equator of the Sun, a phenomenon described by Sporer’s law.

Sunspots usually exist as pairs with opposite magnetic polarity. The magnetic polarity of the leading sunspot alternates every solar cycle, so that it will be a north magnetic pole in one solar cycle and a south magnetic pole in the next. The solar cycle has a great influence on space weather, and is a significant influence on the Earth’s climate since luminosity has a direct relationship with magnetic activity [Lean, et al. 1992]. Solar activity minima tend to be correlated with colder temperatures, and longer than average solar cycles tend to be correlated with hotter temperatures. In the 17th century, the solar cycle appeared to have stopped entirely for several decades; few sunspots were observed during this period. During this era, known as the Maunder minimum or Little Ice Age, Europe experienced unusually cold temperatures [Bonanno., et al. 2008].

Butterfly Diagram

Detailed observations of sunspots have been obtained by the Royal Greenwich observatory since 1874. These observations include information on the sizes and positions of sunspots as well as their numbers. These data show that sunspots do not appear at random over the surface of the sun but are concentrated in two latitudes bands on either side of equator. The great majority of sunspots appear between 5° and 30° north and south of the solar equator. At the start of 11-year sunspot cycles most of the spot lie in 20° to 30° latitudes zone. A butterfly diagram (Fig. 2.5) showing the position of the spots for each rotation of the sun since May 1874 shows that these bands first form at mid latitudes, widen and then move towards equator as each cycle progresses.


Solar Flares

Solar flares are complex transient excitation of the solar atmosphere associated with magnetically active regions of the solar surface evolving enhanced thermal and radio emission. Solar flares are classified on the basis of brightness and area of Hα emission. The solar flare occurrence frequency is highest when the sunspots are more in number and active, and occupy the largest area during their life. The solar flares can generate two types of shock waves (blast waves and driven shocks) in the interplanetary space, which have direct effect on sudden commencements of geomagnetic storms and Forbush decreases in cosmic ray intensity.

Solar flares have been classified as sub flares, 1, 2, 3 and 3+ importances, depending upon the area of Hα emission. The supplementary indices f, n and b are attached to these figures indicating that a flare is faint, normal or bright respectively. The energy released in a single solare flare varies from 1022 joules to 3 × 1023 joules. The frequency of occurrence of solare flares varies, from several per day when the sun is particularly ‘active’ to less than one each week when the sun is ‘quiet’, large flares are less frequent than smaller ones. Solar activity varies with an 11-year cycle. At the peak of cycle there are typically more sunspots on the sun and hence more solar flares. Solar flares extend out from solar corona.

Coronal Holes


Coronal holes are dark regions of the sun with open magnetic field structure in the corona that are less dense and coolar than surrounding aeeas. During the minimum year of the solar cycle they are confined on the sun’s polar regions, while at the solar maximum they can be found at all latitudes. [Antonucci, et al. 2000, Munro, et al. 1972]. Coronal holes were first observed in 1973 when extended observation of the sun was made in X-ray wavelength from Sky-lab. The high speed solar wind originates from the coronal holes [Krieger, et al. 1973]. The high speed solar wind streams are closely connected with the coronal holes. They are situated above the photospheric regions of Uni-polar magnetic field lines open to interplanetary space (Fig. 2.6). The rotation period of the coronal holes is relatively constant with a period of 27 days and with a pole to equator variation of ~3% only. It has been observed that the equatorial coronal holes are usually the extension of the polar holes from either of the solar hemisphere. Moreover, the polar holes are more basic structure near the Sun and the effects of these polar holes are easily communicated to near the Earth interplanetary space.

Coronal Mass Ejections (CMES)

A coronal mass ejection (CMEs) is an ejection of material from solar corona, usually observed with a white light chronograph as is shown in (Fig. 2.8). The ejected material is a plasma consisting primarily of electron and protons, (in addition to small quantities of heavier elements such as Helium, Oxygen and Iron), plus the entrained coronal magnetic field.

The first detection of a CME was made on December 14, 1971 by Tousey (1973) using the 7th orbiting solar observatory (OSO-7). When the ejecta reaches at the earth magnetosphere, compressing it on the dayside and extending the night side tail. When the magnetosphere reconnects on the night side, it creates trillions of watts of power, which directed back towards the earth's upper atmosphere. This process can particularly strong aurora also known as the Northern lights or aurora borealis in (Northern Hemisphere). CME events alongwith solar flare, can disrupt radio transmissions cause power outage (blackouts), and cause damages to satellites electrical transmissions lines. Coronal mass ejections disrupt the flow of the solar wind and produce disturbances that strikes on the earth with some times catastrophic results.


CME's heading straight for the earth are not easy to observe, because are seen against the background of the bright sun. Nevertheless, instrument abroad the solar observatory SOHO, launched in 1996, towards the Langrangion point L1, are able to do so and have been used since early 1997 in space weather prediction. Sometime after the years 2000 NASA hopes to launched a "solar stereo" mission, with a spacecraft in the earth's orbit but 60 or 90 degree away from the earth's position; such a spacecraft would be placed to observed CME's heading from earth from their side. The large angles spectrometric chronograph (LASCO) on the solar and heliospheric observatory (SOHO) has observed a large number of CME's. A typical CME's has three part structure consisting of a cavity of low electron density, a dense core (the prominence which appear as a bright region on chronograph images) embedded in this cavity and a bright leading edge. It should be noted that, however, that many CME's are missing one if these elements or even all three. Most CME's originate from active regions (grouping of sunspots associated with frequent flares). These regions have closed magnetic lines, where the magnetic field strength is large enough to allow the containment of the plasma; the CME's must open these fields lines at least partially escape from the sun. However CME's can also be initiated in the quiet sun regions (although in many cases the quiet regions are recently active). During solar minimum, CME's form primarily in the coronal streamer belt near the solar magnetic equator. During solar maximum, CME's originate from active regions whose latitude distribution is more homogenous. CME range speed is from 20 km/s to 2700 km/s with an average speed (based on SOHO / LASCO measurements between 1996 and 2003) of 489 km/s. The average mass based on chronograph images is 1.6×1015g. Due to the two dimensional nature of the chronograph measureme n t s , these values are lower limits.


The frequency of ejections dependence on the phase of the solar cycle; from about one every other day near solar minimum to 5-6 per day near solar maximum. These values are also lower limits because CME's propagates away from the earth (backside CME's) can usually not be detected by the chronographs. CME's are often associated with the form of solar activity, most notably solar flare, Eruptive prominences and X-ray sigmoid, coronal dimming, EIT and Morton waves, coronal waves, post eruptive arcades.

Signature of CME or ICME

In early studies, interplanetary solar wind data have been used to indirectly infer the nature of solar source of geomagnetic storms. There are in general two kinds of solar sources, CMEs and corotating interaction region (CIRs). The CME counterparts in interplanetary space, conventionally called interplanetary CMEs (ICMEs), can be verified by various solar wind signatures including magnetic clouds [Burlaga, et al. 1981, Klein, et al. 1982] and bidirectional electron fluxes [Gosling, et al. 1987]. Other solar wind features can also be used as CME signatures [Richardson, et al. 1995]. ICMEs are geoeffective because of either the enhancement of an interplanetary magnetic field compressed by CME-driven shocks or the presence of strong magnetic fields carried by CMEs themselves, or both [Bothmer, et al. 1998, Tsurutani, 2001].

The Solar Wind

In previous section we have read plasma come out from the sun, this out flows of plasma from outer atmosphere of Sun is known as solar wind. The temperature of the photosphere is less than 6000K and as a result the sun is most luminous in what we call visible light. The temperature at first falls and then rises rapidly as one move above the photosphere. This behavior is reminiscent of the behavior of the upper atmosphere of the Earth, in which the absorption of solar photons in the stratosphere and the thermosphere create warm layers well above the surface of the Earth. Here too the corona appears to be a region of heating from which heat is conducted downward to the chromosphere and convected outward from the corona by the solar wind.

The solar wind is responsible for the overall shape of Earth's magnetosphere, which is the outermost envelope of the Earth’s environment. Furthermore, the fluctuations in the solar wind speed, its density, direction, and entrained magnetic field strongly affect Earth's local space environment. As during the intense solar activity the levels of ionizing radiation and radio interference varies by factors of hundreds to thousands. The shape and location of the magnetopause and bow shock wave upstream of it can change by several Earth radii, exposing geosynchronous satellites to the direct solar wind. These phenomena are collectively called as “space weather”. Apart from this the solar wind affects the incoming cosmic rays as well as interacting with the atmosphere of other planets. Moreover, planets with a weak or non-existent magnetosphere are subject to atmospheric stripping by the solar wind. The solar wind is divided into two components, the slow solar wind and the fast solar wind respectively. The slow solar wind has a velocity of about 400 km/s, and a temperature of 1.4–1.6×105 K ; a composition that is in close match to the corona. By contrast, the fast solar wind has a typical velocity of 750 km/s, temperature of 8×105 K and it is nearly match the composition of the Sun's photosphere. The slow solar wind is twice as dense as and more variable in intensity than the fast solar wind. The slow wind also has a more complex structure, with turbulent regions and large-scale structures [Kallenrode, et al. 2004, Hassler, et al. 1999, Shrivastava, et al. 2003, 1995].

It has been reported that the slow solar wind appears to originate from a region around the Sun's equatorial belt known as the "streamer belt". The outward coronal flow from these regions i.e. the coronal streamers extend outward from this region, carrying plasma from the interior along with closed magnetic loops [Gruntman, et al. 1999].

Observations of the Sun between 1996 and 2001 showed that emission of the slow solar wind occurred between latitudes of 30–35° around the equator during the solar minimum (the period of lowest solar activity), then expanded toward the poles as the minimum waned. By the time of the solar maximum, the poles were also emitting a slow solar wind [Suess and Steve 1999, Schwenn 2006, Harra, et al. 2008, Rostoker, et al. 1987]. The total number of particles carried away from the Sun by the solar wind is about 1.3 × 1036 per second. Thus, the total mass loss each year is about (2–3)×10−14 solar masses or 6.7 billion tons per hour. This is equivalent to losing a mass equal to the Earth every 150 million years. However, only about 0.01% of the Sun's total mass has been lost through the solar wind.

Interplanetary Magnetic Field (IMF)


Alfven (1954) suggested that a magnetic field embedded in highly conducting plasma tends to behave as if it is 'frozen-in' the plasma. The radially blowing solar wind carrying the frozen-in photospheric magnetic field due to rotation of the sun, twists the field lines which are anchored to the sun, in the form of archimedian spiral [Parker, 1958, Ahluwalia, et al. 1962]. Such a field configuration, which was inferred from cosmic ray flare particle observations, has now been firmly established by direct space measurement. It was shown that the average magnetic field near the earth (outside the magnetosphere) is about 5 RE for the solar field of 1 Gauss. Measurements made on board Mariner 2, Pioneer 5 and IMP-1 had indicated a southward component of field perpendicular to ecliptic. For distances beyond 0.1 AU the radius of the sun may be neglected and the form of the spiral is given by the equation


where, V is the solar wind velocity Ω is the angular velocity of sun and is the heliocentric measured from the reference longitude 0 and r is the radial distance.

The predicted spiral structure of the interplanetary magnetic field was experimentally confirmed first by McCraken (1962) from the cosmic ray solar flare observations. The measurements by instrumentation on board Pioneer 10 and Pioneer 11 spacecraft at the orbit of the Jupiter at a distance of ~0.4 AU have further confirmed the Archimedian spiral configuration of the magnetic field; the garden hose angle may change from 45° to 67°.

Direct measurements have also shown that the magnetic field is distributed in the form of well-defined sector structure. The interplanetary magnetic field seems to be well ordered into sectors with magnetic field being predominately away from the sun (positive) and towards the sun (Fig. 2.9). The existence of sector structure in the IMF, even during the period of maximum solar activity, has been well-established (Wilcox and Colburn, 1969; 1970) [Wilcox, et al. 1969, 1970].



The geomagnetic field originates in the fluid outer core and is generated by electrical currents resulting from convective fluid motions of the highly conducting iron nickel alloy of which the core is made this mechanism is termed the dynamo effect. Even though the fluid motions in the core are quite complex, the magnetic field at the earth's surface (roughly 3000 km away from the core) is nearly as simple as the dipolar magnetic field of simple bar magnet as shown in Fig. 2.10. In addition to sources in the earth's core the magnetic field observable at the earth's surface has sources in the crust and in the ionosphere and magnetosphere. The variations are now made, in the order of low frequency to high frequency variations, in both the space and time domains. The final section describes how the earth's magnetic field of the earth's is often derived as being approximately dipole with field lines emanating from the south geomagnetic pole and converging at the north geomagnetic poles, as depicted in (Fig. 2.10).

Geomagnetic Storm

A geomagnetic storm is a temporary disturbance of the Earth’s magnetosphere caused by a disturbance in the interplanetary medium. A geomagnetic storm is a major component of space weather and provides the input for many other components of space weather. A geomagnetic storm is caused by a solar wind shock wave and/or cloud of magnetic field which interacts with the Earth’s magnetic field. The increase in the solar wind pressure initially compresses the magnetosphere and the solar wind magnetic field will interact with the Earth’s magnetic field and transfer an increased amount of energy into the magnetosphere. Both interactions cause an increase in movement of plasma through the magnetosphere (driven by increased electric field inside the magnetosphere) and an increase in electric current in the magnetosphere create magnetic force which pushes out the boundary between the magnetosphere and the solar wind. The disturbance in the interplanetary medium which drives the geomagnetic storm may be due to a solar coronal mass ejection (CME) or a high speed stream (co-rotating interaction region or CIR) of the solar wind originating from a region of weak magnetic field on the Sun’s surface. The frequency of geomagnetic storms is increases and decreases with the sunspot cycle. CME drive storms are more common during the maximum of the solar cycle and CIR driven storms are more common during the minimum of the solar cycle.

There are several space weather phenomena which tend to be associated with a geomagnetic storm or are caused by a geomagnetic storm. These include: Solar Energetic Particle (SEP) events, geo-magnetically incluced current (GIC), ionospheric disturbances which cause radio and radar scintillation, disruption of navigation by magnetic compass and auroral displays at much lower latitudes than normal. In 1989, a geomagnetic storm energized ground induced currents, which disrupted electric power distribution throughout most of Quebec province and caused aurora as far south as Texas.

In 1931, Chapman and others wrote in their article, a New Theory of Magnetic Storms, which sought to explain the phenomenon of geomagnetic storms [Chapman, et al. 1930]. They argued that whenever the Sun emits a solar flare it will also emit a plasma cloud. This plasma will travel at a velocity such that it reaches Earth within 113 days. The cloud will then compress the Earth’s magnetic field and thus increase this magnetic field at the Earth’s surface [Ferraro, 1933]. A geomagnetic storm is defined [Gonzalez, et al. 1994] by changes in the Dst [Tsurutani, et al. 2003] (disturbance-storm time) index. The Dst index estimates the globally averaged change of the horizontal component of the Earth’s magnetic field at the magnetic equator based on measurements from a few magnetometer stations. Dst is computed once per hour and reported in near-real-time. During quiet times, Dst is between +20 and -20 nano-Tesla (nT).

A geomagnetic storm has three phases [Gosling J.T., et al. 1987]: an initial phase, a main phase and a recovery phase. The initial phase is characterized by Dst (or its one-minute component SYM-H) increasing by 20 to 50 nT in tens of minutes. The initial phase is also referred to as a storm sudden commencement (SSC). However, not all geomagnetic storms have an initial phase and not all sudden increase in Dst or SYM-H are followed by a geomagnetic storm. The main phase of a geomagnetic storm is defined by Dst decreasing to less than -50 nT. The selection of -50 and approximately -600 nT. During of the main phase is typically between 2 and 8 hours. The recovery phase is the period when Dst changes from its minimum value to its quiet time value. The period of the recovery phase may be as 8 hours or as long as 7 days.



Unlike magnetic storms, substorms are magnetic disturbances occurring in a limited region of the earth, near the geomagnetic poles. The most obvious effect of a substorm is the occurrence of auroral displays near the north and south magnetic poles (Fig. 2.11). As was discovered in the last century, aurora are most frequent not at the magnetic poles themselves but in almost circular zones some 4000 km in diameter centred on the magnetic poles, known as the auroral ovals. The visible emission seen in an aurora occurs 90-130 km above the earth’s surface, and is due to the excitation of atoms in the atmosphere by energetic electrons. The characteristic red and green colours of strong auroral displays arise from spectral lines of neutral oxygen atoms in the earth’s atmosphere, with wavelengths of 630.0 nm (red) and 557.7 nm (green).

Interaction of IMF with Earth Magnetosphere

The solar wind, carrying with it the interplanetary magnetic field, flows out from the sun into the solar system and so encounters planets and other bodies orbiting the sun. How does the wind interact with them, particularly those planets like the earth that have magnetic field of their own?

Measurements from spacecraft and the ground have given us a very detailed picture in the case of the earth, one which can be broadly applied to the other magnetized planets that have been explored by spacecraft in recent years. What happens is that, as the solar wind moves through the lines of the earth’s field, electric currents are induced that alter the field’s configuration, compressing it on the sunward side and causing the wind flow to be diverted, avoiding a region surrounding the earth. The result is a cavity, known as a magnetosphere, having a boundary (the magnetopause) that is nearly spherical on the sunward side but drawn out into a long cylinder, the magnetotail, on the night side (see Fig.2.12). Expressed in units of the earth’s radius (1 RE=6371 km), the distance of the magnetopause from the earth along the sunward direction is 10 RE and the magnetotail’s diameter is 50 RE. Outside the magnetopause, the solar wind continues to flow, but within it the terrestrial magnetic field dominates. The magnetosphere is filled with plasma having a large range of densities and temperatures; the charged particles constituting the plasma, unlike the particles making up the other regions of the earth’s atmosphere, are bound by electromagnetic, i.e. non-gravitational, forces.


Near the earth’s surface, the terrestrial magnetic field is dipole-like, i.e. like that produced by a bar magnet near the earth’s core, and is tilted by about 20º to the earth’s rotation axis. On the sunward side of the magnetosphere, the terrestrial field lines are closed, as are those on the night side out to between 8 and 15 RE from the earth. The density of the plasma filling this, the inner magnetosphere, decreases sharply above a level corresponding to the magnetic field lines that are 4-5 RE distant from the earth along the sun-earth line. Inside this level, the plasmapause, is a region called the plasma-sphere where the gas has a density of 107-108 particles/m3 and originates from the ionosphere, a region of the earth’s atmosphere that is partly ionized by the sun’s radiation. Outside the plasma-pause the density is only106 particles/m3. Much of this plasma, too, is ionospheric, but particles are present with much higher energies that come from the solar wind, particularly via the Polar Regions, or cusps; the reason is that some field lines at the cusps connect with regions very near the magnetopause, allowing the solar wind particles to leak into the magnetosphere. In the magnetotail beyond about 8-15 RE, the field lines are open in their normal configuration and parallel to the earth-sun direction. The magnetotail has a northern lobe with field lines pointing towards earth, a southern lobe with field lines pointing away from earth, and an intermediate region called the plasma sheet where oppositely directed field lines lie close to each other forming a cross-tail current, in the form of a sheet. The plasma sheet eventually converges to a neutral line some 100 RE from the earth. It is the site of various plasma processes giving rise to geomagnetic and auroral phenomena. Because of the tilt of the earth’s magnetic axis to its rotation axis and the tilt of the rotation axis relative to the orbital plane, the plasma sheet is generally north or south of the ecliptic plane by up to 4 RE.

The magnetopause is actually a layer with a thickness of several hundred kilometers rather than a sharp boundary to the magnetosphere. There is slight motion in it according to the state of the solar wind, and in particular there is a rippling motion along the magnetotail boundary. Ahead, or ‘upstream’, of the magnetopause by about 3 RE on the sunward side is a standing shock front, known as the bow shock, whose presence is due to the fact that, relative to the solar-wind flow, the earth and its magnetosphere are moving faster than either the sound or Alfven speed. Its properties, like the magnetopause, depend on the solar wind, in particular the interplanetary magnetic field carried with it: when the direction of this field is almost parallel to the shock, the shock front is thin and well defined, but when nearly perpendicular it is rather wide and ill-defined. Some solar-wind particles incident on the shock are scattered back into a sunward direction. Between the bow shock and the nose of the magnetopause, a region known as the magnetosheath, the solar-wind flow continues at reduced speed, and is somewhat compressed and irregular.

There is a convection of plasma in the magneto tail as a combined result of a strong electric field set up across the magneto tail and magnetic field there. The convicted plasma moves towards the earth, then round it, leaving the magnetosphere on its sunward side. Some of the electron and ions in this convicted plasma become trapped in the field of the inner magnetosphere, following helical paths along the field lines and mirroring near the terrestrial magnetic poles. The motion between one mirror point and the other typically takes a fraction of a second. Owing to the decrease of field strength going away from the earth, there is a drift motion, with electrons drifting eastwards, ions westwards. This separation of charge sets up a current – the- ring current – flowing round the earth. The direction of an electric current is conventionally in the opposite direction to the flow of electrons, so the ring current flows westwards.

Measurement of Geomagnetic Storm

Presently, magnetic storm and sub-storm is monitored at ground level by observing changes in the Earth’s magnetic field over periods of seconds to days, by observing the surface of the Sun and by observing radio noise created in the Sun’s atmosphere. The Sunspot Number (SSN) is the number of sunspots on the Sun’s photosphere in visible light on the side of the Sun visible to an Earth observer. The number and total area of sunspots are related to the brightness of the Sun in the extreme ultraviolet (EUV) and X-ray portions of the solar spectrum and to solar activity such as solar flares and coronal mass ejections (CMEs). For the ground base observation we use varies geo-magnetic indices which calculate the change in earth magnetic field at varies part of earth. Some most common indices describe below :


Dst index is an estimate of the magnetic field change at the Earth’s magnetic equator due to a ring of electrical current at and just earthward of GEO. The index is based on data from ground-based magnetic observatories between 21º and 33º magnetic latitude during a one hour period. Stations closer to the magnetic equator are not used due to ionosphere effects. The Dst index is compiled and archived by the World Data Center for Geomagnetism, Kyoto.

Kp and ap index

Kp/ap index: ‘a’ is an index created from the geomagnetic disturbance at one mid-latitude (40° to 50° latitude) geomagnetic observatory during a 3 hour period. ‘K’ is the quasi-logarithmic counter-part of the ‘a’ index. Kp and ap are the average of K and a over 13 geomagnetic observatories to represent planetary-wide geomagnetic disturbances. The Kp/ap index indicates both geomagnetic storms and substorms (auroral disturbance). Kp/ap is available from 1932 onward.

AE Index

AE index is compiled from geomagnetic disturbances at 12 geomagnetic observations in and near the auroral zones and is recorded at 1 minute intervals. The AE index is made public with a delay of two to three days, which severely limits its utility for space weather applications. The AE index indicates the intensity of geomagnetic sub storms except during a major geomagnetic storm when the auroral zones expand equator ward from the observatories.


  • • Abhyankar, K.D. (1977), Bull. Astr. Soc. India, 5, p. 40.
  • • Adams, F.C.; G. Laughlin and G.J.M. Graves (2004), Red Dwarfs and the End of the Main Sequence, Revista Mexicana de Astronomia y Astrofisica, 22, p. 46.
  • • Ahluwalia, H.S. and A. J. Dessler (1962), Planetary Space Sci., 9, p. 195.
  • • Antonucci, E.; M. Dodero and S. Giordano (2000), Solar Phys., 197, pp. 115–134.
  • • S. Basu, S.; W.J. Chaplin; Y. Elsworth; R. New and A.M. Serenelli (2009), Astrophys.J., 699, 1403.
  • • Basu, S. and H.M. Antia (2008), Physics Reports 457(5-6), p. 217.
  • • Bessell, M.S.; F. Castelli and B. Plez (1998), Astronomy and Astrophysics, 333, p. 231.
  • • Bonanno, A.; H. Schlatt and L. Paterno (2008), Astronomy and Astrophysics, 390(3), p. 1115.
  • • Bothmer, V. and R. Schwenn (1998), Geophys Ann., 16, p. 1.
  • • Burlaga, L.F.; E. Sittler; F. Mariani and R. Schwenn (1981), Geophys J. Res., 86, p. 6673.
  • • Burton, W.B. (1986), Space Science Reviews, 43, p. 244.
  • • Chapman, S. and V.C.A. Ferraro (1930), Nature, 129, p. 129.
  • • Emslie, A.G. and J.A. Miller (2003), Particle Acceleration, Dynamic Sun, Cambridge University Press, p. 275.
  • • Erdelyi, R. and I. Ballai (2007), Astron. Nachr, 328(8), p. 726.
  • • Ferraro, V.C.A. (1933), The Observatory, 56, p. 253.
  • • García, R.A.; S.T.-Chièze; S.J.J.-Reyes; J. Ballot; P.L. Pallé; A. Eff-Darwich; S. Mathur and J. Provost (2007), Science, 316, p. 1591.
  • • Gonzalez, W.D.; J.A. Joselyn; Y. Kamide; H.W. Kroehl; G. Rostoker; B.T. Tsurutani and V.M. Vasyliunas (1994), Geophys J. Res., 99(A4), p. 5771.
  • • Gosling, J.T.; D.N. Baker; S.J. Bame; W.C. Feldman; R.D. Zwickl and E.J. Smith (1987), Geophys J. Res., 92, p. 8519.
  • • Gruntman, M.A. (1999), Geophys J. Res., 99, p. 19213.
  • • Hansteen, V.H. and E. Leer (1997), Astrophysical J., 482(1), p. 498.
  • • Harra, L.; R. Milligan and B. Fleck (2008), Hinode: Source of the slow solar wind and superhot flares, ESA.
  • • Hassler, D.M.; I.E. Dammasch; L. Philippe; B. Pål; W. Curdt; E.M. Helen; J.-C. Vial and K. Wilhelm (1999), Science, 283, p. 810.
  • • Kallenrode, M.B. (2004), J. Geophys. Res., 103, p. 6917.
  • • Klein, L.W. and L.F. Burlaga (1982), J. Geophysical Res., 87, p. 613.
  • • Kogut, A.; C. Lineweaver; G. F. Smoot; C.L. Bennett; A. Banday; N.W. Boggess; E.S. Cheng; G. de Amici; D.J. Fixsen; G. Hinshaw; P.D. Jackson; M. Janssen; P. Keegstra; K. Loewenstein; P. Lubin; J.C. Mather; L. Tenorio; R. Weiss; D.T. Wilkinson and E.L. Wright (1993), Astrophysical J., 419, p. 1.
  • • Krieger, A.S.; A.F. Timothy and E.C. Roelof (1973), Solar Phys., 29, p. 505.
  • • Lada, C.J. (2006), Astrophysical J. Letters, 640, p. 1.
  • • Lean, J.; A. Skumanich and O. White (1992), Geophysical Res. Lett., 19(15), p. 1591.
  • • Lodders, K. (2003), Astrophysical J., 591(2), p. 1220.
  • • Mc-Cracken, K.G. (1962), J. Geophys. Res., 67(2), p. 423.
  • • Munro, R.H. and G.L. Withbroe (1972), Astrophys. J., 176, p. 511.
  • • Parker, E.N. (1958), Astrophysics J., 128, p. 664.
  • • Richardson, I.G. and H.V. Cane (1995), J. Geophysical Res., 100, p. 397.
  • • Riley P.; J.A. Linker and Z. Mikic (2002), J. Geophysical Res., 107(A7), p. SSH 8-1.
  • • Rostoker, G.; S.I. Akasfu; W. Baumjohann and Y. Kamide (1987), Space Sci. Rev., 46, p. 93.
  • • Schwenn, R. (2006), Solar Physics, 3, p. 1.
  • • Shrivastava, P.K. and R.P. Shukla (1995), Solar Physics, 154, p. 177.
  • • Shrivastava, P.K. and K.L. Jaiswal (2003), Solar Phys., 214, p. 195.
  • • Solanki, S.K. and T. Ayres (1994), Science, 263, p. 64.
  • • Suess, S. (1999), Overview and Current Knowledge of the Solar Wind and the Corona, The Solar Probe, NASA/Marshall Space Flight Center.
  • • Tsurutani, B.T. (2001), in Space Storms and Space Weather Hazards, ed. I.A. Daglis (Dordrecht: Kluwer), p. 103.
  • • Tsurutani, B.T.; W.D. Gonzalez; G.S. Lakhina and S. Alex (2003), J. Geophys. Res., 108 (A7), p. 1268.
  • • Wilcox, J.M. and D.S. Colburn (1969), J. Geophys. Res., 74, p. 2388.
  • • Wilcox, J.M. and D.S. Colburn (1970), J. Geophys. Res., 75, p. 63-66.
  • • Woolfson, M. (2000), Astronomy and Geophysics, 41, p. 1.