Chapter 1. Introduction about Sun

Transcription

1 Introduction about Sun

2 1.1. Introduction Sun is the only star and most important feature in our solar system. Large and violent eruptions are occurring on the Sun due to surface and coronal activities like., sunspots, coronal mass ejections (CMEs), solar X-ray flares, faculae, prominences and coronal holes. The Sun is with diameter of about 1,392,000 km (about 109 times that of Earth) and mass about kg (330,000 times that of Earth). The other important characteristics are the following : (1) Sun is about 99.86% of the total mass of our solar system, (2) Chemically, about 75% of the solar mass consists of hydrogen, while the rest is mostly helium, (3) Less than 2% consists of heavier elements, including carbon, nitrogen, oxygen, neon, magnesium, silicon and iron and (4) The mean distance of the Sun from the Earth is approximately km (known as 1 Astronomical Unit (1 AU) and it is a unit of length in Astrophysics). Sun is almost spherical and it consists of mostly hot hydrogen plasma. This plasma is tenuous and gaseous near the surface, but gets denser down towards the Sun's core. Since the Sun is composed of gaseous plasma, the rotation of the Sun varies with latitude. The regions of the Sun near its equator rotate once in every 25 days. The Sun's rotation rate decreases with increasing latitude, so that its rotation rate is slowest near its poles. At its poles the Sun rotates once in every 36 days. The structure and behaviour of the Sun is determined by the mass, momentum and energy conservation laws and the mode of energy transport. The Sun is an oblate spheroid, like all the major bodies in the solar system. However in a first approximation, in some of the studies (Mestel, 1965; Mack and Robbins, 1972; Townsend, Owocki and Ud-Doula, 2007), the effects of rotation and magnetic fields may be neglected so that the Sun is taken to be spherically symmetrical. The internal structure of the Sun is divided into three different regions called core, radiative and convective zones. The temperature of the core is 15.6 million kelvin and the pressure is 340 billion atmospheres. At the center of the core the 1

3 Sun's density is more than 150 times that of water. The core is the only region in the Sun that produces an appreciable amount of thermal energy through nuclear fusion. The rest of the star is heated by energy that is transferred outward from the core to the outside layers. Once energy is produced in the core of the Sun, it needs a way to travel from the solar center to the outer regions. The physical transport of energy from its production site to the surrounding regions can be done in a number of ways. However, the most efficient means of transferring energy near the core is by radiation. Consequently, the region surrounding the core of the Sun is known as the radiation zone and this is a region of highly ionized gas and the energy transport is primarily by photon diffusion. In the radiation zone of the Sun the temperature is little cooler than the core and as a result some atoms are able to remain intact. These intact atoms are able to absorb energy, store it for a while, and then later emit that energy as new radiation. In this manner the energy that is generated in the core is passed from atom to atom through the radiation zone. It takes over 170 thousand years for the energy released in the core of the Sun to get out of the radiation zone. Outside the radiation zone it requires a new transport mechanism to continue its journey to the surface. This new method of transport is required because outside the radiation zone the temperature is relatively cool (2 MK), whereas the temperature of the radiation zone is nearly 5 MK. At this temperature the atoms will absorb energy, but because things are cool and dense the atoms do not release it so readily. Consequently the transfer of energy by radiation slows down significantly. The most efficient means of energy transfer is now convection and this region of the Sun's interior is known as the convection zone. The hotter material near the top of the radiation zone (the bottom of the convection zone) rises up and the cooler material sinks to the bottom. As the hot material reaches the top of the convection zone it begins to cool and sink, and as it sinks it heats up again and will rise. This produces a rolling motion much like a convection process in a pot of boiling water. The hot material follows a direct path through the convection 2

4 zone and the energy is transferred much faster than the rate of energy transfer from radiation zone to convection zone. It takes only a little more than a week for the hot material to carry its energy to the top of the convection zone. Figure 1 shows the internal structure, solar atmosphere and various coronal features of the Sun as clipped from Solar and Heliospheric Observatory (SOHO) by European Space Agency and National Aeronautics and Space Administration (ESA and NASA). Figure 1 Various basic features of the Sun clipped from SOHO images. 3

5 1.2. Solar Atmosphere In this thesis, it is attempted to study the activities seen on the surface and atmosphere of the Sun. The Solar atmosphere is interesting and important for its characteristics as the features of the solar atmosphere has importance in Sun-Earth interactions and affect the space weather of the Earth. Solar atmosphere is classified into three regions based on the heliocentric distance from the surface of the Sun. They are (1) the visible layer of the Sun called photosphere, where sunspots, coronal holes and faculae occur, (2) Chromosphere, a region just few thousand kilometers above the photosphere, where flares and plages occur and this region is seen clearly during solar eclipse times and (3) the outermost solar corona extending more than one million kilometers from surface of the Sun where prominences and coronal mass ejections (CMEs) occur Photosphere The outer visible layer of the Sun is called the photosphere with a temperature of 5, 800 K and a density of about kg m -3. The temperature of the photosphere of the Sun ranges from 4,500 K to 6,000 K (with an average temperature of 5800 K). Features like sunspots, faculae and granules can be observed in the photosphere with a simple telescope (along with a good filter to reduce the intensity of sunlight to safe levels). We can also measure the flow of material in the photosphere using the Doppler Effect. These measurements reveal additional features such as super-granules as well as large scale flows and a pattern of waves and oscillations. 4

6 Figure 2 A typical image of the Photosphere (visible layer). The photosphere is composed of convection cells called granules which are cells of gas each approximately 1000 km in diameter with hot rising gas in the center and cooler gas falling in the narrow spaces between them. Each granule has a lifespan of only about eight minutes, resulting in a continually shifting "boiling" pattern. Granules may group into super-granules up to 30,000 km in diameter with lifespan of up to 24 hours can also be seen. The photosphere is effectively the radial zone at which the density of ionized hydrogen atoms becomes low enough that the photons can escape unhindered. It is only a few hundred kilometers thick. The granulated, orange-peel like, appearance of the Sun of the photosphere is shown in Figure 2. This image was taken using TRACE (Transition Region and Coronal Explorer) instrument in Stanford-Lockheed Institute for Space Research by NASA. 5

7 Chromosphere The chromosphere (known as "color sphere") is a thin layer of the Sun roughly, 2,000 km just above the photosphere. The chromosphere is more visually transparent than the photosphere. Due to its low density, it is relatively transparent, resulting in the photosphere being regarded as the visual surface of the Sun. It is red in color and the spectrum of the chromosphere is dominated by the deep red H α spectral line of hydrogen. This may be seen directly with the naked eye only during a total solar eclipse. The structure of the chromosphere is as shown In Figure 3. This is a spectrum of the solar chromosphere during the solar eclipse of March 7, For reasons not fully understood, the temperature of the chromosphere is hotter than that of the photosphere. Figure 3 The typical image of chromosphere (reddish layer region). 6

8 The photosphere is closer to the center of the Sun and its temperature is around 5800 K but the chromosphere temperature varies from 4500 K to as high as 20,000 K. The acoustic turbulence may be the source of this higher temperature, as the result of the dispersion of magneto-hydrodynamic waves over the solar surface. The most common solar feature within the chromosphere are spicules, long thin fingers of luminous gas which appear like the blades of a huge field of fiery grass growing upwards from the photosphere below. Spicules rise to the top of the chromosphere and then sink back down again over the course of about 10 minutes Solar Corona There is a region around the Sun, called corona, extending more than one million kilometers from chromosphere, where the temperature is about 2 MK and hence it emits X-ray radiations. This is the region where the solar wind originates. The corona can be seen only during solar eclipses, when the main radiation from the photosphere is blocked by the passage of the Moon. Artificial solar eclipses can be obtained by occulting the photosphere of the Sun by a disc in front of the objective of a telescope (coronagraph) and images of corona may be obtained. Due to its low density of particles (roughly g cm -3 ) we can see right through it, and the underlying chromosphere, down to the photosphere. The temperature of the core of the Sun is about 15 MK and it drops to 5800 K at the surface of the photosphere. However the temperature of the corona is found to be nearly 2 MK. This hot temperature at the corona requires a permanent heating mechanism, or the plasma would cool down in about an hour. There are many mechanisms which could heat some gas above the surface of the Sun, but none of those mechanisms could account for the large rate of heating necessary to heat the corona to these temperatures. This phenomenon remained a mystery for more than 50 years. 7

9 The magnetic field is playing an active role in heating the solar corona. The heat energy is coming from the active regions, where the spectacular giant loops are seen in Ultra-Violet (UV) and X-rays and these giant loops are found to disappear during solar minima. The heating of the corona is linked to the interaction of the magnetic field lines. Because the laws of electromagnetism prohibit the intersection of two magnetic field lines, every time magnetic field lines come close for crossing they are said to be "rearranged" and this is called magnetic reconnection and it is expected to heat the solar corona continuously. It is a fairly inefficient source of energy, but the sheer number of these small magnetic patches on the surface of the Sun makes the process a viable solution to the 50 year old problem for heating. The approximate length of time these patches remain active is not certain. As a direct consequence of this theory, the heating process should occur much closer to the surface of the Sun than previously thought, but no one really knows how close (Schmelz, 2003) Solar Activity The Sun emits radiation in a wide range of the electromagnetic spectrum from long radio waves to x-rays, including high-energy particles. The transient phenomena occurring in the solar atmosphere are as follows: (1) sunspots and faculae in the photosphere, (2) flares and plages in the chromosphere and (3) prominences and coronal mass ejections in the corona and all these can be grouped together under the term solar activities. These are the most important solar surface activities to disturb the space weather of the Earth and hence are found to be important in Sun-Earth connections Sunspots The sunspots are dark regions on the photosphere, because they are at lower temperature with high magnetic field strength region than the surrounding regions of the photosphere. They are known to possess intense localized magnetic fields (of 8

10 the order of 0.4 T). In a typical solar spot, the umbra is the inner, darker and cooler (3400 K) region of a sunspot compared to the surroundings and width is about 20,000 km. In a typical solar spot, the penumbra is the outer, relatively less darker compared to an umbra but more darker compared to the surrounding photosphere region. It is shaped like an annulus (a ring) surrounding the darker and cooler umbra. A pore is a smaller sunspot but without a penumbra. The sizes of pores are about 2,500 km wide and are less darker than an umbra of a typical sunspot. Granules are regions of the Sun where hot solar material comes up to the solar surface. Granules are about 1,000 km across and only exist for about 5 to 10 minutes before they fade away. Granulation covers the visible surface (the photosphere) of the Sun and is associated with large scale fluid motions at and just below the photosphere. The brighter and central regions correspond to rising of hotter plasma; and the darker and narrow lanes correspond to sinking of colder plasma. Typical speeds in granular flows are of the order of a few km s -1. These regions are clearly shown in Figure 4. Figure 4 The different parts of the sunspots during the 23 rd solar cycle maximum. 9

11 The number of sunspots on the surface of the Sun varies with respect to time periodically and it is called as solar cycle and also known as the sunspot cycle. It is found that the solar magnetic field varies in a 22-year cycle. If one could stand on the surface of the Sun with a compass, it would be pointing towards one pole of the Sun for 11 years, and then it would switch around and point at the other pole for 11 years. From the early records of sunspots it was seen that the Sun went through a period of inactivity in the late 17 th century. Very few sunspots were seen on the Sun from about AD 1645 to AD Although the observations were not as extensive as in later years, the Sun was in fact well observed during this time and this lack of sunspots was well documented. This period of solar inactivity was found to coincide to a climatic period called the "Little Ice Age" when rivers that were normally ice-free were found to froze and snow fields remained year-round at lower altitudes. Figure 5 The yearly averaged sunspots cycle. 10

12 At present, Sun-Earth relations specifically space weather predictions of the Earth are related to transient solar activities and these investigations form a thrust area. Observations of solar spots 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 latitude bands on either side of the equator. In general, the position of the spots for each rotation of the Sun shows that these bands first form at mid-latitudes, widen, and then move toward the equator as each cycle progresses. The yearly averaged sunspots number during the period is shown in Figure 5. During the solar cycle, on the surface of the photosphere, there are active regions, quiet Sun regions and coronal holes. Active regions are located in areas of strong magnetic field, visible as sunspots groups in optical wavelengths or from magnetograms. Whenever the number of sunspots is maximum, the Sun is called active Sun and when number of sunspots is minimum the Sun is named as quiet Sun. Sunspots occur when the Sun s local magnetic field in a region is stronger than average magnetic field of the surroundings in the photosphere. Sunspots appear in pairs, connecting a loop of the Sun s magnetic field lines at two foot points. When gas travels along these looped field lines, one gets an arch-shaped solar prominence. The comparison of the Sun at solar minimum (right of July 2008 in Figure 6) and at solar maximum (left of August 2002 in Figure 6) as seen in extreme ultraviolet light from SOHO (Solar and Heliospheric Observatory) is shown in Figure 6. Generally the northern and southern polar zones of the Sun have been found to be darker than the equatorial zones. These regions are called as coronal holes. Sometimes coronal holes are found to be present in the regions other than polar zones also. Today it is fairly clear that these zones are dominated by open magnetic field lines, which act as efficient conduit for flushing heated plasma from the corona into the solar wind. Because of this efficient transport mechanism, coronal holes are empty of plasma most of the time, and thus appear much darker than the 11

13 surroundings, where heated plasma up flowing from the chromosphere remains trapped until it cools down and precipitates back to the chromosphere. The heliographic position of active regions is typically confined within latitudes of ±40 ο from the solar equator. Figure 7 shows the soft X-ray image of the extended solar corona recorded on 1992 August 26 by the Soft X-ray Telescope (SXT) with many active regions and coronal holes. Sunspot groups typically exhibit a strongly concentrated leading magnetic polarity, followed by a more fragmented trailing group of opposite polarity. Because of this bipolar nature, active regions are mainly made up of closed magnetic field lines. Due to the permanent magnetic activity in terms of magnetic flux emergence, flux cancellation, magnetic reconfigurations, and magnetic reconnection processes, a number of dynamic processes such as plasma heating, flares, and coronal mass ejections occur in active regions. Figure 6 Comparison of active (left) and quiet (right) Sun regions (23 rd solar cycle). 12

14 Figure 7 The active region with coronal loops and coronal hole Coronal Mass Ejection (CME) Coronal Mass Ejection (CME) is an ejection of material from the solar corona, usually observed with a white-light coronagraph. The ejected material is highly magnetized hydrogen plasma consisting of electrons and protons. In addition, small quantities of heavier elements are present such as helium, oxygen, and iron. The magnetic reconnection is responsible for CME and solar flares. Magnetic reconnection is the name given to the rearrangement of magnetic field lines when two oppositely directed magnetic fields are brought together. This rearrangement is accompanied with a sudden release of energy stored in the oppositely directed fields. On the Sun, magnetic reconnection may happen on solar arcades, called a series of closely occurring loops of magnetic lines of force. These lines of force quickly reconnect into a low arcade of loops, leaving a helix of magnetic field unconnected to the rest of the arcade. The sudden release of energy in this 13

15 reconnection causes the solar flare. The unconnected magnetic helical field and the material that it contains may violently expand outwards forming a CME. Generally, the CMEs are observed by the white light coronagraphs. The Solar and Heliospheric Observatory (SOHO) mission's white light coronagraphs observed nearly coronal mass ejections (CMEs) between 1996 and The measured properties of all these CMEs are documented in an on-line catalog ( The Large Angle and Spectrometric Coronagraph (LASCO) (Brueckner et al. 1995) on board the SOlar and Heliospheric Observatory (SOHO) provide unprecedented capabilities for direct detection of CMEs at heights from 2 Ro to 32 Ro (Ro = km). The coronagraph, LASCO C2 covers the heliocentric distance range of 2-6 R o and LASCO C3 covers the heliocentric distance range of R o. The innermost coronagraph C1 operated only for the first 2.5 years; therefore we will not include C1 observations here. Both C2 and C3 have the same image size ( pixels) with a pixel size of 11.2 and 56.0 arc seconds, respectively. The SOHO/LASCO team use standard LASCO software available as IDL routines in SOLARSOFT (Freeland and Handy, 1998) to run movies of LASCO images and measure the increase in height of CMEs as they expand away from the Sun. They typically use the running difference movies to better identify frame-toframe changes in the corona. LASCO is able to take images of CMEs at different times and the corresponding different heights of the solar corona by blocking the light coming directly from the Sun with an occulting disk, creating an artificial eclipse within the instrument. The position of the solar disk is indicated as a white circle with the structure of a CME in Figure 8. From this SOHO/LASCO coronagraph, the various parameters of CMEs such as width, position angle and mean speed etc are measured. From the height-time measurement, the linear mean speed and mean residual acceleration (herein after mean acceleration) of the CMEs are also estimated and they are listed in the on-line catalog. 14

16 Figure 8 The structure of a CME. A typical CME has a three part structure consisting of a cavity of low electron density, a dense core embedded in this cavity, and a bright leading edge (Hundhausen et al. 1984). It should be noted, that many CMEs are missing any one of these parts, or even all the three. The current knowledge of CME kinematics indicates that the CME starts with an initial pre-acceleration phase characterized by the slow rising motion, followed by a period of rapid acceleration away from the Sun until a near constant velocity is reached. Some balloon CMEs lack this threestage evolution. The white - light observations of coronal mass ejections (CMEs) often show the classic "three-part" structure. Figure 8 shows the SOHO/LASCO C2 coronagraph image which has the three part structure of the CMEs on 24 th November The width, position angle, cosmic ray hits and coronal streamers of the CME are also marked in this figure. 15

17 Figure 9 The height-time graph of CME taken during 24 th November From the LASCO images one can easily find the width, central position angle (CPA) of corresponding CMEs. Figure 9 shows height-time measurement of a typical CME by using the SOHO/LASCO C2 and C3 coronagraph. In this figure asterisks and square symbols (data points) represent the LASCO C2 and C3 observation. The few corresponding LASCO C2 and C3 running difference images are inserted in this figure. The acceleration and kinetic energy of the CMEs are determined and listed in SOHO/LASCO CME on-line catalog. If the CMEs with high magnetized plasma passes near the Earth, its electromagnetic field will distort the Earth's magnetic field, the energetic particles of CME will funnel through the Earth's polar regions into the ionosphere, and the heat due to CME will expand the Earth's atmosphere. These effects can cause dramatic problems to power grids, communications and navigation systems, satellite and the space weather at our Earth. 16

18 Solar soft X-ray flares The solar flares are another important transient features to affect the geospace weather. A solar flare is a sudden brightening observed over the surface or the limb of the Sun. A flare occurs when magnetic energy that has buildup in the solar atmosphere is suddenly released. This is interpreted as a large energy release of up to J (about one sixth of the total energy output of the Sun in every second). A flare process is associated with a rapid energy release in the solar corona, believed to be driven by stored non-potential magnetic energy and triggered by instability in the magnetic configuration. Such an energy release process results in acceleration of non-thermal particles and in heating of coronal/chromospheric plasma. These processes emit radiation in almost all wavelengths such as: radio, white light, EUV, soft X-rays, hard X-rays during large flares. Solar flares affect all layers of the solar atmosphere (photosphere, chromosphere, and corona). When the plasma is heated to tens of millions of kelvins, then the particles like electrons, protons, and heavier ions are accelerated to very high speed near to the speed of light. A typical solar flare is shown in Figure 10. Flares occur in active regions around sunspots, where intense magnetic fields penetrate the photosphere to link the corona to the solar interior. The same energy releases may produce coronal mass ejections (CMEs), although the relation between CMEs and flares is still not well established. X-rays and UV radiation emitted by solar flares can affect Earth's ionosphere and disrupt long-range radio communications. Direct radio emission at decimeter wavelengths may disturb operation of radars and other devices operating at these frequencies. The frequency of occurrence of solar flares varies from several per day when the Sun is said to be "active" to less than one every week when the Sun is said to be "quiet", following the 11-year cycle (the solar cycle). Solar flare along with a CME is potentially dangerous to spacecraft and especially to people in space. 17

19 To obtain the information on the surface events (flares and eruptive filaments or prominences) that may be related to CMEs, we have taken flares data from the GOES satellite (Geostationary Operational Environmental Satellite) maintained by NGDC (National Geophysical Data Center). A typical GOES soft X-ray flare profile is observed by low and high (1 8 Å and Å, respectively) energy channels represented by two solid line shown in Figure 11. According to NGDC, the starting time of a X-ray flare is defined as the four consecutive one-minute soft X-ray fluxes in the GOES 1 to 8 Å band that meet the following criteria: (i) all 4 values are above the B1 threshold (flux greater than 10-7 W m 2 ), (ii) all 4 values are strictly increasing, and (iii) the last value is greater than 1.4 times the value which occurred 3 min earlier. Physically, this definition corresponds to the starting of the impulsive phase. In this study, we use the starting, ending, peak times, location and peak flux (Importance) values of the flare in soft X-rays reported by NGDC. The starting time is usually several minutes prior to the starting of their associated hard X-ray emission or the starting of their associated H α flare. Figure 10 The solar flare is taken by the Solar Dynamics Observatory (SDO) on 13 February

20 Figure 11 The GOES soft X-ray flare profile taken on 2 nd April The X-ray flare intensity is measured in units of power per area. The X-ray flare importance or peak fluxes are classified into A, B, C, M and X when the intensity is < 10-7, , , and > 10-4 W m -2 respectively. Any numerical value after the class refers to the multiplicity factor. To determine the exact intensity of the flare one has to multiply the number in the X-ray classification of that flare by the value of its class listed above. For example, a C5.9 flare would have an intensity of 5.9 x 10-6 W m Solar Radio Bursts Solar radio bursts occur when non-thermal electrons are accelerated during the solar eruption process. These non-thermal emissions represent the dynamic processes occurring in the corona. Non-thermal radio emission does not have the characteristic signature curves of blackbody radiation. In fact, it is quiet the opposite, with emission increasing at longer wavelengths. The most common form of non-thermal emission found in astrophysics is called synchrotron emission. It arises by the acceleration of charged particles within a magnetic field. Most 19

21 commonly, the charged particles are electrons since they are less massive compared to protons. As the energetic electrons encounter a magnetic field, they spiral around it. Since the spiral is continuously changing direction, it is in effect accelerating and emitting radiation as shown in Figure 12. The frequency of emission is directly related to speed of the electron. A wider spiral implies emission at comparatively longer wavelengths (continuum emission). It is important to note that, unlike thermal emission, synchrotron emission is polarized. As the emitting electron is viewed side-on in its spiral motion, it appears to move back-and-forth in straight lines. Wild and McCready (1950) classified solar radio bursts into five major types. A summary of present classification, including burst characteristics and associated solar phenomena, is given in Table 1. While various bursts types can occur as single events, they often occur in combination. In other words, it is quiet common to see a complex radio event originating from a single solar outburst. This is illustrated in Figure 13, which presents a typical dynamic spectrum following a large flare. This diagram illustrates each major type of solar radio burst in a typical configuration following a large flare. It should be noted that not all of these features are observed in every flare. Figure 12 Schematic diagram shows the example of Synchrotron emission. 20

22 A detailed description of the five major types of radio bursts and their significance are given below: (i) Type-I events are categorized as short and narrowband events that usually occur in great numbers together with a broad-band continuum, which typically last for hours or days. They are thought to be associated with solar flare and possibly eruptive prominence (Kai, Melrose, and Suzuki, 1985), both of which can result in geomagnetic and ionospheric storms. Figure 13 Classification of radio bursts. (ii) Type-II emissions appear to be emission strips slowly drifting from high frequency (300 MHz) to low frequency (10 MHz) in the dynamic spectrum (emitting frequency Vs time). They often have double structures, exhibiting fundamental and second harmonic emissions. Their frequency drift results from the outward motion of a magneto-hydrodynamic shock wave at the leading edge of material ejected during a flare and/or CME. From a statistical study of type-ii bursts Mann et al. (1996) found that their drift rate is 0.16 ± 0.11 Hz s -1 and their life time is in the 21

23 range 5 to 15 min. The drift rate can be measured for estimating coronal shock speeds if coronal density information is given. (iii) Type-III bursts are narrow-band ones that sweep rapidly from high to low frequencies. They often occur in groups and, on occasions, in storms with an accompanying continuum. Like type-ii bursts, they can exhibit dual structures, with fundamental and second harmonic emissions. The bursts are believed to originate in plasma oscillations associated with the ejection of an electron jet. In general, these bursts are indicative of the presence of an active region on the solar disk. More importantly, intense type-iii groups are often produced in the initial stage of large flares. (iv) Type-IV bursts are defined as a smooth continuum of broad-band bursts primarily in the meter range ( MHz). They have been divided into a variety of different sub-classes by various authors (Robinson, 1985b; Pick, 1986). It is also noted that some of these sub-classes correspond to different physical phenomena. For example, there is said to be a strong association between the incidence of stationary type-iv bursts and the emission of protons (Robinson, 1985b), but no such association with moving type-iv bursts (Kai, 1979). These, and other subclasses (flare continua), can all appear quite similar. Indeed, they are often only distinguishable by subtle differences in their fine structure or circular polarization. These bursts are associated with some major flare events beginning 10 to 20 min after the flare maximum, and can last for hours. (v) Type-V bursts are broad-band continua which often follow type-iii bursts or groups. They are short-lived, typically lasting for only a few minutes. Such events sometimes appear to merge with their accompanying, or subsequent, bursts making individual features difficult to discern. The emission is believed to be produced by synchrotron emission from the fast-rising electron jet. 22

24 Table 1 Classification of solar radio bursts Type Characteristics Duration Frequency Range (MHz) Associated Phenomena I Short Narrow-band Continuum Hrs-days Active regions Prominence Flares II Slow drift Second harmonics 5-30 min Flare/CMEs Proton emission Shock wave III Fast drift Group Second harmonics Hrs-days Active regions Flares Stationary: Broad-band Continuum Hrs-days Flares Proton emission IV Moving: Broad-band Slow drift 30 min 2 hrs Prominence Shock wave Flare Continua: Broad-band Smooth continua 3-45 min Flares Proton emission V Smooth Short-lived Follow type-iii 1-3 min Active regions Flares 23

25 However, out of five radio bursts the type-ii radio bursts is most important among all other radio bursts. Type-II bursts were first identified by Payne-Scott, Yabsley, and Bolton, 1947, who recognized the importance of mass motion for these bursts. Wild and McCready (1950) classified them as type-ii bursts to distinguish from the fast-drifting type-iii bursts. In a dynamic radio spectrum, a type II bursts appears as an emission band slowly drifting from high to low frequencies (McLean, 1967; Smerd, Sheridan, and Stewart, 1975; Nelson and Melrose, 1985; Zlotnik et al. 1998). On the basis of the slow frequency drifts, these bursts have been interpreted as the radio signature of coronal shock waves which can be generated by flares and/or CMEs (Robert, 1959; Smerd, Sheridan, and Stewart, 1975; Wild, Sheridan, and Trent, 1959). The metric type-ii bursts (hereinafter m-type-ii bursts) have shock speeds in the range of km s -1 (Robinson, 1985a). During the last decade there has been a controversy about the two shock scenario. Gopalswamy et al. (1998) and Reiner at al. (2001) proposed that a flare-associated coronal shock, detected by metric type-ii radio bursts and this will decay to an MHD (Magneto hydrodynamic) wave close to the Sun. According to the above classification, type-ii radio bursts appear to be emission stripes slowly drifting from high to low frequencies in the dynamic spectrum. It has been interpreted that type-ii bursts are the signatures of shock waves associated with solar flares and/or coronal mass ejections (Gopalswamy et al. 1998; Cliver, Webb, and Howard, 1999). As the shock propagates from the Sun, the density of the ambient elections decreases and so does the excited frequency (f = fp) where fp is the plasma frequency of the ambient medium defined by (1/2π) (n e 2 /m ε) where n, e, m and ε represent electron density, charge, effective mass of electron and permittivity of the medium. The drift rate can be converted into shock speed by employing an appropriate coronal electron density model (Newkirk, 1961; Smerd, Sheridan, and Stewart, 1975; Robinson, 1985a). 24

26 Figure 14 Schematic dynamic spectrum showing different type-ii bursts confined to various wavelength ranges. Figure 15 Type-II burst components in meter (m), decameter -hectometer (DH) and kilometer (km) wavelength domains. 25

27 From the ground based and space-borne radio spectrographs, three different type-ii radio bursts are observed: (i) meter wavelength, (ii) deca-hectometer wavelength and (iii) kilometer wavelength type-ii radio bursts. The coronal type-ii radio bursts are in meter to kilometer wavelength and are assumed to be generated by the Langmuir turbulence excited by electrons that are accelerated at the MHD shocks created by flares and/or CMEs (Nelson and Melrose, 1985). The type-ii bursts occurring in the decameter-to-hectometer wavelength range (DH-type-II bursts) are caused by CMEs produced in the interplanetary space (Gopalswamy et al. 2001). Radio observations in the 14 1 MHz frequency band are continuously recorded by the Wind/WAVES instrument (Bougeret et al. 1995) in the Wind spacecraft (Acuña et al. 1995). The appearance of type II bursts in various wavelength domains is shown in Figure 14: (1) bursts confined to the metric (m) domain; (2) bursts starting in the m domain but continuing into the DH domain; (3) bursts confined to the DH domain; (4) bursts starting in the DH domain and continuing into the kilometer (km) domain; (5) bursts having counterparts in all the wavelength domains, m-to-km; (6) bursts confined to the km domain. In the schematic picture, we have not shown the details of harmonic structure or bandsplitting, which may or may not be present in all the events. In the above item number (6), the components in various spectral domains may and may not have direct continuity. The m- type-ii burst discussed in item number (1) occurs typically above the ionospheric cutoff at 20 MHz, observed from ground based radio telescopes. Radio emission from the Sun at longer decameter and kilometer wavelengths cannot penetrate the terrestrial ionosphere, so space borne instruments observe type-ii bursts discussed in item numbers 2-6. Coronal densities similar to the ionospheric densities occur at a heliocentric distance of 3 R ο, which is also considered to be the location of the source surface of the solar magnetic field. The ambient medium beyond the source surface is considered to be the interplanetary (IP) space. Thus, the bursts at frequencies above the ionospheric cutoff are known as coronal type-ii bursts (m-type-ii bursts), while 26

28 the ones occurring at frequencies below the cutoff are known as the IP (DH, km) type-ii bursts (Gopalswamy, 2004). Occasionally, one can observe population discussed in item number 2 (m-to-dh) bursts using ground based instruments at geographical locations where the ionospheric cut-off is well below the nominal 20 MHz (Erickson, 1997). One has to combine ground and space based observations to see the continuation from the m-domain to DH and km domains. Type-II radio bursts components in metric (m), decameter - hectometric (DH) and kilometric (km) wavelength domains combined from ground based spectrograph and space borne instrument (Wind/WAVES) are shown in Figure 15 (Gopalswamy, 2011) Solar-Terrestrial relationships The Sun is the important source of geo-space weather consequences. During the last few decades we have known the hazards due to the Sun-Earth interactions. The large and violent eruptions of plasma and magnetic fields from the Sun, known as CMEs, are the origin of geomagnetic storms (National Academy of Sciences [NAS], 2008). The flare associated CME shock waves create solar energetic particles (SEPs), which are high-energy particles consisting of electrons and solar wind ions (mainly protons). The CMEs occurring on the front side of the Sun are the potential source of geomagnetic storms because they can directly impact Earth s magnetosphere with high kinetic energy. It takes approximately two to three days after a CME launched from the Sun for a geomagnetic storm to reach Earth and to affect the Earth s geomagnetic field (NERC, 1990). They can inject large amount of mass and magnetic fields into the heliosphere, causing major geo-magnetic storms and interplanetary shocks, a key source of solar energetic particles (SEPs). The geoeffective CME depends on the existence of southward component of the magnetic field in the sheath and/or ejecta portions. Particularly CME may be dangerous for astronauts and upset satellite electronics. The Earth is well protected by its magnetic field. But when the large flare associated CME comes with its own magnetic structure, there is some chance that the two field points may be in 27

29 opposite directions so that Earth and CME attract each other like two magnets. The Earth s magnetic shield is partially canceled by the magnetized plasma from the CME s field, and vast number of protons and electrons find easy access to our atmosphere and extending the auroras far beyond the polar circles, this is named as geoeffective storms. The study of these worldwide disturbances of Earth s magnetic field are important in understanding the dynamics of solar-terrestrial environment and furthermore because such storms can cause life threatening power outrages, satellite damage, communication failure and navigational problems (Joselyn and McIntosh, 1981; Lakhina, 1994; Gonzalez et al. 1994). Auroras are the natural light display in the sky particularly in the high latitude (Arctic aurora- australis and Antarctic aurora- borealis) regions, caused by the collision of energetic charged particles with atoms in the high altitude atmosphere (thermosphere). The charged particles originate in the magnetosphere and in the solar wind and are directed by the Earth's magnetic field into the atmosphere. In addition, the changing the magnetic field may induce electric currents in the Earth s surface or in power transmission line or power grid. A famous example is the 1989 Quebec blackout (Lerner, 1995) Geomagnetic storms Geomagnetic disturbance are generally represented by geomagnetic storms and sudden ionosphere disturbances (SIDs). These are caused by the disturbances originated at solar atmosphere, interplanetary (IP) shocks and / or stream interfaces associated with high speed solar wind streams (Shriastava and Agarwal, 1990; Kuznetsov et al. 1998). The geomagnetic storm is defined by changes in the Dst (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. The Dst is computed once per hour and reported in near-real-time. The size of a geomagnetic storm is classified as moderate ( -50 nt >minimum of Dst > -100 nt), intense (

30 nt > minimum Dst > -250 nt) or super-storm ( minimum of Dst < -250 nt). When Dst < -50 nt the storm is considered to be significant (Tsurutani et al. 1988; Webb et al. 2000). CMEs also can result in increasing significant levels of solar energetic particles (SEPs) at Earth (Reames, 1999). When the > 10 MeV proton intensity exceeds 10 particle flux units (particle flux unit; 1 pfu = 1 particles.cm -2 s -1 sr -1 ) the particle events are significant and the associated CMEs are considered SEP effective. The strength of IMF (Interplanetary Magnetic Field) and its fluctuations have also shown to be most important parameter affecting the geomagnetic field condition. South direction of IMF, allows sufficient energy transfer from the solar wind into the Earth magnetosphere through magnetic reconnection (Howard et al. 1985; Zhang and Burlaga, 1988; Kaushik and Shriastava, 1999). The geomagnetic storms are the major component of space weather and provide 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 pressure of the solar wind 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 as shown in Figure 16. Both interactions cause an increase in movement of plasma through the magnetosphere (driven by increased electric fields inside the magnetosphere) and an increase in electric current in the magnetosphere and ionosphere. A geomagnetic storm comprises characteristic field changes in time sequential phases usually conveniently described as two or three phases, they are: (i) initial (not necessarily in all storms), (ii) main, and (iii) recovery phases as shown in Figure 17. The initial phase begins with a storm sudden commencement (SSC) and it is distinguished as a positive excursion in the ground magnetic field. As the CME hits the magneto-pause, the magneto-pause is compressed and then increases the ground field at the equator. The main phase is due to an increase in energetic ions and electrons in the inner magnetosphere, where they become trapped on closed 29

31 magnetic field lines and drift around the Earth, thus creating the ring current. Following this main phase, it is lasting from half an hour to several hours. At this time, Dst index reaches negative values of hundreds of nt and also ring current is increased. The strong cross-tail electric field pushes the plasma-pause closer to the Earth, and peels off the outer layers of the plasma sphere. The large, radially outward (in equatorial plane) pointing electric field builds up at the earthward edge of the plasma sheet, which is the region of closed field lines in the equatorial magneto tail. The ring current begins to decay during the recovery phase. It is lasting from tens of hours to a week and Dst index gradually returns to the normal level. The ring current ions gradually decrease. The expanding plasma sphere brings cold ionospheric plasma in connection with the ring current. Figure 16 A illustration of the large flare associated CME produced geomagnetic storm produced with geo-magnetosphere. 30

32 Figure 17 An illustration of Dst index data with three phase of geomagnetic storm. The number of geomagnetic storms increases and decreases with respect to solar activity (Solar cycle). CME driven storms are more common during the maximum of the solar cycle and CIR (Co-rotating Interaction Region) 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 geomagnetic storms. Based on all the above discussions the CMEs and solar soft X-ray flares and their associated lower coronal type-ii radio bursts activities are the most important features of the Sun to study the various consequences of geospace weather. There is a long time controversy of the origin of meter wavelength type-ii radio bursts in the vicinity of the Sun. However, the type-ii bursts occurring in the decameter-to-hectometer wavelength range (DH-type-II bursts) are found to be driven by CMEs (Kaiser et al. 1998; Reiner et al. 2000; Gopalswamy et al. 2001). The DH wavelength domain (DH-type-II) bridges the gap between the traditional metric band of ground based radio telescopes and kilometric band of space-borne radio instruments (Gopalswamy et al. 1998; Kaiser et al. 1998; Reiner and Kaiser, 1999). Using this DH domain type-ii radio data, one has to find the shock strength 31

33 i.e., whether a single shock was formed in the metric domain and it was extended to DH and kilometric domain. On the other hand, whether two different shocks are formed in the two different wavelength domains or not? The work in this thesis attempted to find answers for the above. Some more interesting features found in the analysis of type-ii bursts, flares and CMEs are also discussed in the following chapters of this thesis. The organization of the work reported in this thesis is given in the next section Thesis organization In the Chapter 2, the relationship between type-ii radio bursts appearing in the meter (m) and decameter-to-hectometer (DH) wavelength ranges are analyzed. The associated X-ray flares and coronal mass ejections (CMEs) are also reported. The selected sample of events are divided into two classes using the drift plots of type-ii radio bursts pairs (log frequency- time): (i) Class I, representing those events where DH-type-II bursts are not continuation of m-type-ii bursts and (ii) Class II, where the DH-type-II bursts are extensions of m-type-ii bursts. Our study consists of three steps: (i) comparison of characteristics of the Class I and II events; (ii) correlation of m-type-ii and DH-type-II burst characteristics with X-ray flare properties and (iii) correlation of m-type-ii and DH-type-II burst characteristics with CME properties. In Chapter 3 the investigations are carried out on: (i) the statistical properties of m- and DH-type-II bursts; (ii) time lags between onsets of flares and CMEs associated with type-ii bursts; and (iii) statistical properties and relation between flares and CMEs of Class I and Class II events. In this chapter, the significant differences between the properties of m- and DH-type-II bursts of Class I and Class II events are found. 32

34 In Chapter 4, the kinematic properties of this Class I and Class II radio-loud CMEs and associated with flares and type-ii radio bursts pairs are discussed. In addition, the importance of Class I and Class II radio-loud CMEs their associated various Sun-Earth connection are also discussed. In this chapter, the following investigations are carried out: (i) the height is estimated at which the CME attained their peak speed within the LASCO field of view and at which the corresponding maximum speed of the CMEs reported, (ii) CME nose height is estimated at the start time of DH-type-II bursts, and (iii) the initial acceleration of the CMEs are estimated. In addition, it is also checked in this chapter whether flare associated Class I and Class II radio-loud CMEs are the good particle accelerator or not? The purpose of this Chapter 5 is to address the following questions: (i) whether the acceleration of CME is more important to generate shocks in the interplanetary medium and (ii) among the three different groups of CMEs (accelerated/decelerated/constant speed), what kinds of differences in the characteristics of associated activities (DH-type-II bursts and flares) exist? In this chapter, it is also studied the characteristics of DH-type-II radio bursts, flares and CMEs based on the classification of the CME acceleration within the LASCO field of view. A sample of 212 DH-type-II associated with CMEs is classified into three populations: (i) Group I (43 events) representing DH-type-II associated CMEs are accelerating in the LASCO field view (a > 15 m s 2 ); (ii) Group II (99 events) representing approximately constant velocity CMEs ( 15 < a <15 m s 2 ) and (iii) Group III (70 events) representing decelerating CMEs (a < 15 m s 2 ). Our study consists of the following three steps: (i) statistical properties of DH-type-II bursts of Group I, II and III events; (ii) analysis of time lags between onsets of flares and CMEs associated with DH-type-II bursts and (iii) statistical properties of flares and CMEs of Group I, II and III events. The results of this detailed analysis are discussed in this chapter. 33