The Sun, at the center of our solar system, is the source of life on the Earth. Its

Transcription

1 Chapter 1 THE HELIOSPHERE 1.1 Introduction The Sun, at the center of our solar system, is the source of life on the Earth. Its variations, however subtle, may contribute to changes in our global environment. It provides models, available for study, of astrophysics processes occurring throughout the universe. The outer atmosphere of the Sun expands into space to form the solar wind. In turn, The solar wind in turn carves out a region of interstellar space, several times the orbit of the most distant planet in dimension, known as the heliosphere. The heliosphere is the region of space that is dominated by the solar magnetic field, that is drawn out into space by the solar wind, a bubble in space. It defines the volume of space over which our Solar influence predominates. That influence reaches beyond the way what we normally think of as the solar system. As the strength and level of activity of the solar wind varies in space and time, the size and behavior of the heliosphere too varies (figure 1.1). The closest boundary of the heliosphere is thought about 100 AU from the Sun. The solar wind has its initial acceleration up to 40 solar radii and maintains its speed of about 400 km 8-1 up to a radial distance 7

2 Chapter 1: THE HELIOSPHERE 8.../, /....'... M~~biC Fi;~d Lines.... ~."'..."'.....,.. /.L I~.... Figure 1.1: The heliosphere of about 70 AU. Beyond this distance its density decreases because of its expansion and hence the flow slows down to subsonic speed. The solar wind then continues as a subsonic flow until the pressure of the interstellar gas becomes larger than the combined pressure of the solar wind and the frozen-in magnetic field. This location is the heliopause, the boundary of the heliosphere, which is beyond 100 AU. Inside the heliosphere, the solar wind plasma flows radially outward from the Sun, carrying with it the interplanetary magnetic field (IMF) which is spiral in shape, with about 20 full windings. As the Sun and solar system travel about the galaxy, it carries the heliosphere along with it. But at the same time the interstellar medium moves, but in a different direction. The combined motion of the Sun and the interstellar wind compress the heliosphere in the direction of incoming interstellar wind and elongates the heliosphere in downwards. The heliosphere is filled with particles and fields of solar origin, connected with

3 Chapter 1: THE HELIOSPHERE 9 solar transients like the solar wind, solar flares, and coronal mass ejections. But, the Sun is not the only source of particles that flow through the interplanetary space. Other sources include the comets, the planets, the interstellar medium and the galaxy. These heliospheric particle population cover a large dynamic range in intensity and in energy. The heliosphere acts as a shield, protecting us from the bulk of cosmic rays generated far away in the universe which would otherwise bombard the Earth continuously, causing damage to living cells. Without the heliosphere, life would certainly have evolved in a different way. Sometimes life would not have even emerged. 1.2 Solar wind Origin of solar wind The Sun is a massive, luminous ball of gas. It consists of about 90 % hydrogen (by number) and 10 % helium with a small fraction of heavier elements, such as carbon, oxygen, and iron. Within a region of about 0.3 solar radii from the center of the Sun, the temperature and density are high enough ( 10 7 K and 100 g/cm 3 ) to create fusion of protons into helium nuclei. The energy released from this process diffuses outward as "hard" X-rays which are degraded into radiation of longer wavelengths by continuous absorption and emission of the photons by the gas surrounding the. core. Above a distance of about 0.7 solar radii from the center, the energy is carried away by convection. As the hot gas rises, it becomes less dense until it is finally transparent, and the transported energy is radiated out into interplanetary space. The radiation occurs is the photosphere which is regarded as the visible surface of the Sun. Above the photosphere lies a layer of transparent gas called the chromosphere. The chromosphere is much dimmer than the photosphere which is visible seen only with the aid of a coronagraph or during a total solar eclipse. It appears as a pearly

4 Chapter 1: THE HELIOSPHERE 10 white structure on eclipse photographs, parts of which extend to several solar radii from the edge of the Sun. The coronal temperature ranges from 1 to 2 X 10 6 K, and it is therefore in a highly ionized plasma state. For an example, the spectral lines from iron with up to 12 electrons removed have been observed, and the lightest elements are fully ionized. The corona is a higwy structured region of the solar plasma. This structure is imposed by the solar magnetic field which extends from the solar surface out into the corona. Since the corona is plasma (i.e., a collection of positively charged nuclei and negatively charged electrons), it is an excellent electrical conductor. As a result of this conductivity, the coronal plasma can move along the magnetic field lines; but it cannot across magnetic field lines.there are two types of magnetic field lines are there; closed and open. Closed field lines are anchored at two points in the photosphere which extend into the corona as in the case of magnetic field loops seen in the solar prominence. Open field lines are anchored at only one point in the photosphere and extend into interplanetary space. It is in these open field regions, the corona can expand outward in the form of the solar wind. On the surface of the Sun, the magnetic field is compressed in to small flux tubes. As we move up through the solar atmosphere, at some regions, the gas pressure decreases in proportional to the magnetic pressure. These magnetic flux tubes, rooted in the surface of the Sun, control the structure of the corona. Thermally driven solar wind cannot produce the observed high-speed solar streams and the mechanism of supplying huge amount of energy for the supersonic flow of solar wind from the coronal holes. The attempts to relate theory with remote observations or ecliptic spacecraft and telescope observations nearly in 1 AU hardly satisfy any standard model. Probably the acceleration process will be better understood if in situ observations of the inner heliosphere are made possible within few solar radii of the Sun. Parker's hydrodynamical solar wind model (Parker, 1958) is a valid approximation of the solar wind

5 Chapter 1: THE HELIOSPHERE 11 in quiet conditions. A comparison of the model results with the observations reveals that the hydrodynamic model is more appropriate to describe the slow wind originating in the streamer belt with its complex magnetic field structures than the fast solar wind blowing directly out of the coronal holes, a situation which is much closer to the assumptions inherent in the model. Observations suggest that the solar wind and its sources might be even more complicated. In particular, it appears that the solar wind does not expand radially. An expansion factor can be defined as the ratio between the cross sections of the flux tube at the source surface and in the interplanetary space. High speed solar wind streams from coronal holes are connected with small expansion factors while low latitude slow streams are linked with large expansion factors. These slow streams originate at the boundaries of the coronal holes, their strong expansion causes them to fill the space above the low latitude closed magnetic structures. Since the expansion factor can be inferred from the magnetic field strength and distribution on the source surface, magnetograms can be used to predict the solar wind speed at the orbit of the earth (Wang and Sheeley, 1994) Interplanetary medium The expanding coronal gas or solar wind fills the interplanetary space. The solar magnetic fields, embedded in the coronal plasma, are carried into space by the solar wind to form the interplanetary magnetic field (IMP). Beyond some solar radii, the solar magnetic field is dominated by the solar wind flow which expands almost radially away from the Sun. Because of the solar rotation, the point where the open field line is anchored to the Sun moves and as a result, the interplanetary magnetic field has the form of a spiral shape. At the orbit of the Earth, IMP makes an angle of about 45 to the radial direction. Further out, the field is nearly transverse to the radial direction. At 1 AU the average speed of the solar wind is about 400 km S-1.

6 Chapter 1: THE HELIOSPHERE 12 This speed is by no means constant. The solar wind can reach speeds in excess of 900 km S-l and it can travel as slowly as 300 km 8-1. The average density of the solar wind at 1 AU is about 7 protons/cm 3 with large variations. The solar wind confines the geomagnetic field and governs the phenomena such as geomagnetic storms and aurorae. The solar wind confines the magnetic fields of other planets as well. As the solar wind expands, its density decreases as the inverse square of its distance from the Sun. At the heliopause, the solar wind can no longer "push back" the fields and particles of the local interstellar medium. The solar wind slows down from 400 km S-l to perhaps 20 km S-l. It exhibits a complex spatial and temporal structure on many different scales. The solar wind can be associated with the structure on the Sun itself. In the outer heliosphere, all the solar wind plasma streams in the equatorial plane become compressed at least once by the time it reaches a distance about 20 AU from the Sun. The basic structure of the solar wind in the equatorial plane in the distant heliosphere thus differs considerably from that observed at 1 AU. The solar wind which has its initial acceleration up to 40 solar radii, maintains its free flow speed of about 400 km 8-1 at a radial distance of about 70 AU. Beyond this distance, its density decreases because of its expansion. Hence, the solar wind flow slows down to subsonic speed. The solar wind continues as a subsonic flow until the pressure of the interstellar gas becomes larger than the combined pressure of the solar wind and the IMF. The interaction of solar wind with interstellar gas is highly variable, mainly because of the relative motion of the interstellar gas due to galactic movement of the solar system and the ever changing solar wind momentum flux (Lazarus et al., 1995). During the solar minimum, the slow solar wind originates from regions close to the current sheet at the heliomagnetic equator. During the solar maximum, the solar wind originates above the active regions in the streamer belt. Compared to the fast solar wind, the slow wind is highly variable and turbulent, often containing many

7 Chapter 1: THE HELIOSPHERE 13 large scale structures such as magnetic clouds. Intermittent increase in solar wind is effected by powerful coronal mass ejections where the solar wind speed is in the range 400 km S-1 to 2000 km S-I. Three dimensional plasma measurements made by the Ulysses spacecraft showed a pronounced latitudinal variation of the solar wind. The solar wind increases from about 450 km S-1 in the equatorial plane to 800 kin S-1 above the poles. The transition from high to low speed solar wind is found to be abrupt. The high speed wind from polar coronal hole is steady and the low speed wind occurring between ± 20 of the equator, is erratic and highly variable. As the solar wind flows outward from the Sun, the density rapidly drops with distance as r- 2 and the collisional mean free path becomes much larger than it was in the corona. The solar wind contains an equal number of electrons and protons, along with a few heavier ions, and it escapes out continuously from the surface of the Sun. This wind leads to a mass loss of about 10 million tons of material per year. Though this mass looks like a large quantity, it is insignificant in comparison with the total mass of the Sun. The solar wind escapes primarily through coronal holes dominated by open field lines, which are found predominantly near the polar region. In the equatorial plane, the magnetic field lines are more likely to close on themselves, particularly during the periods of low solar activity. These closed field lines trap the hot coronal gases, leading to enhanced X-ray emissions from these hotter regions and decrease the over all flow of the solar wind. The fast wind has a flow speed between 400 km S-1 and 800 kin S-1 and the average density of 3 ions/cm 3 at 1 AU. The slow wind moves with a speed between 300 and 400 km S-1 and the average density is about 8 ions/cm 3 at 1 AU.

8 Chapter 1: THE HELIOSPHERE Temperature of the solar wind The theoretical expectation of the dependence of the solar wind temperature is neither simple nor clear. Indeed, the measured electron temperature at 1 AU, Te = 1-3 X 10 5 K, which is typically several times greater than the measured proton temperature Tp = x 10 5 K, indicating that the two fluid gases are not coupled thermally. Theoretical projection of the variation of temperature with radial distance is closer to the measured solar wind electron temperature. The proton temperature measured by Pioneer 10 and 11, as the function of radial distance, shows some what faster rate of decrease than that of electrons. The solar wind temperature decreases away from the Sun with a 1/R x dependence out to about 40 AU. Solar wind temperature in the outer heliosphere is much higher than it would be expected for simple adiabatic expansion. This suggests the existence of a source of heating in the outer heliosphere. The possible way for a heat source is the conversion of bulk kinetic energy in to thermal energy by the interaction between high and low-speed streams. The overall temperature found to increase from 30 AU to 50 AU, then starts to decrease again from 50 AU to 63 AU and again starts to increase. Smith and Marsden (1995) suggested that the increase in temperature outside AU is probably due to transfer of energy from pickup protons to thermal protons. In general, the temperature structure of the solar wind is determined by many factors like adiabatic cooling, stream-interaction, pickup proton heating and variation in the initial temperature at the solar wind source region. The average solar wind temperature appears to vary with the solar cycle. It is lower near the solar maximum and higher near the solar minimum. The solar cycle variation of the temperature in the outer heliosphere is much smaller than that observed at 1 AU. The variation of the temperature with heliolatitude can be deduced from the solar stream structure. During solar minimum, the strongest high speed streams observed at high latitude has the

9 Chapter 1: THE HELIOSPHERE 15 highest temperature. The proton temperature in the high latitude heliosphere is well described by a simple radial variation with little latitudinal effect Composition of the solar wind The solar wind consists primarily of electrons and protons; a-particles and other ionic species are present at very low abundance levels. At the orbit ofearth, 1 AU from the Sun, solar wind parameters observed are as shown in the Table 1.1. Interplanetary scintillation (IPS) study of solar wind plasma shows strong latitudinal dependence and the variation is more predominant in solar minimum than solar maximum. Proton density Electron density He2+ density Flow speed (nearly radial) Proton temperature Magnetic field 6.6 cm cm cm kin S x10 5 K 7 x 1O- 9 tesla Table 1.1: Solar wind parameters near 1 AU The table 1.1 shows the main constituents of solar plasma; protons and electrons constitute 95 %, and the rest consists of mainly doubly ionized helium nuclei and heavy ion species. The helium abundance is highly variable during energetic transients. Flux measurements of both light and heavy ion species have been made in the solar wind by many spacecrafts. Bame et al. (1976) reported that the average He to H density ratio is (NlIe/N H ) = The ratio of H to He has been measured to be as low as and as high as 0.25 during flare-associated solar wind. The hydrogen is entirely in the form of protons and helium in the form of a-particles, the

10 Chapter 1: THE HELIOSPHERE 16 average oxygen to helium ratio is No/N He = 0.01, and the ratio of silicon to oxygen and iron to oxygen are NsdNo = 0.21 and NFe/No = The presence of higwy charged states of oxygen, silicon, and iron, reflect the thermal conditions in the inner corona. The Ulysses experiments provide information about the latitudinal distribution of solar particles in solar wind (Feldman et al., 1998). The a-particle to proton abundance showed a consistently low value 4.4 ± 0.6% in the high speed solar wind in the polar region. This is comparable to those measured in the high speed streams near ecliptic. If at all there is a variation, that is confined to the presence of slow or variable solar wind in the ecliptic region. Ulysses also found an increased presence of a-particle abundance in CME's near the ecliptic. CME's at high latitudes showed no enhancement of a-particle abundance ratio. The momentum flux of the solar wind was found to be larger over the polar regions than in the equatorial region (Phillips et al., 1995c; Goldstein et al., 1995) Solar and Interplanetary magnetic field The Sun has a strong and complex magnetic field, and much of the solar activity appears to be directly connected with the properties of the magnetic field. Sunspots are places where very intense magnetic lines of force break through the solar surface. Magnetic field lines loop through the solar atmosphere and interior to form a complicated web of magnetic structures. Many of these structures are visible in the chromosphere and the corona. The sunspot cycle results from the recycling of magnetic fields by the flow of material in the interior. The prominence seen floating above the surface of the Sun are supported, with magnetic fields. The streamers and loops seen in the corona are shaped by magnetic fields. Magnetic fields are virtually the root of all the features we see on and above the Sun. The magnetic field of the Sun is measured in the photosphere. The magnetic field of the photosphere can

11 Chapter 1: THE HELIOSPHERE 17 be probed precisely by observing the Zeeman splitting (Lin et al., 2000). The solar magnetic activity is the consequence of the dynamo at the bottom of the solar convection zone. Dynamo action changes the solar poloidal field (associated with dipole) to a toroidal field (associated with sunspot magnetic field). These two magnetic fields are alternately destroyed and created in a solar cycle which lasts for nearly 22 years (Hale cycle). This may account for the oscillatory nature of the solar cycle (toroidal field oscillation). The solar dipole field has been measured and studied extensively using magnetograms. Ulysses vector helium magnetometer/fluxgate instruments provide the first direct in situ measurement of the poloidal and toroidal solar magnetic fields. The solar magnetic field changes greatly during the maximum phase of the solar cycle, as do the coronal and heliospheric magnetic fields. One of the striking features of the solar activity cycle is the reversal of polarity of the polar magnetic field. The solar dipole magnetic field reverses its polarity around the maximum phase of the solar activity. The characteristics of the polar reversal play a significant role in determining the coronal and interplanetary magnetic fields. The solar magnetic dipole is nearly aligned with the solar rotation axis around the solar minimum. The orientation of the solar magnetic dipole axis with respect to the solar rotational axis is quite variable. The solar dipole is most noticeably tilted relative to the rotation axis during the declining phase of solar activity. The solar magnetic field is not dipolar near the solar activity maximum. Coronal magnetic field The hot corona is also permeated by magnetic field, which is quite strong in some locations (thousands of Gauss in active regions above sunspots). Magnetism dominates the structure and dynamics of the solar corona. Coronal magnetic field is responsible for a wide range of phenomena from being the carrier of magneto hydro dynamic

12 Chapter 1: THE HELIOSPHERE 18 (MHD) waves to heat the corona, producing the gyro-synchrotron radiation in the radio wavelength range. Under the tnfiuence of this magnetic field, the electrons gyrate to produce cyclotron frequency. The radio emission at microwave wavelengths by the gyrated electrons is used to determine the coronal magnetic field. The polarization dependence of the emitted radio waves provides a rough estimation of the magnetic field strength of the corona. These observations are more easily made above or near sunspots. The average radio magnetic field strength to the optical one shows that the decrease of the magnetic field of a sunspot from the photosphere to the boundary of corona is only about 20 %. The spectral and polarization observations of local sources made with RATAN-600 using high resolution in the wavelength range cm, allow one to measure the maximum magnetic fields of the corresponding sunspots at the height of the chromosphere-corona transition region (CCTR). This method (Akhmedov et ai., 1982) is based on determining the short wavelength limit of gyroresonance emission of the LS and relating it to the third harmonic of gyrofrequency. Gary and Hurford (1987) used Owens valley frequency agile interferometer to measure a field strength of 1400 Gauss, and Akhmedov et ai. (1986) used the RATAN 600 to measure a field strength of 1700 gauss in the corona. Despite the importance of magnetic field in the physics of the corona and the tremendous progress made recently in the remote sensing of solar magnetic fields, reliable measurements of the coronal magnetic field strength and orientation do not exist. This is largely due to the weak coronal magnetic fields and the difficulty associated with observing extremely faint coronal emission. Magnetic field strength varies as a function of height for the active regions, the potentiality of the magnetic field configuration (i.e., whether it is potential or force-free) can be determined and directly compared with the extrapolation results. Using a very sensitive infrared spectropolarimeter to observe the strong near-infrared coronal emission line Fe above active regions, Lin et ai. (2000) have suc-

13 Chapter 1: THE HELIOSPHERE 19 ceeded in measuring the weak Stokes V circular polarization profiles resulting from the longitudinal Zeeman effect of the magnetic field of the solar corona. From these measurements, they inferred field strengths of 10 and 33 G from two active regions at two heights. Dulk et al. (2004) made observations of three flares made at 5 and 15 GHz with the VLA and derived 360 to 660 G as the maximum field strength in flaring loops. Interplanetary magnetic field The interplanetary magnetic field (IMF) is a part of the coronal magnetic field that is carried into interplanetary space by the solar wind. In addition to being a very good thermal conductor, the solar wind plasma is an excellent electrical conductor. The electrical conductivity of the plasma is so high that the solar magnetic field is frozen into the solar wind flow, as it expands away from the Sun. While one end of the interplanetary magnetic field (IMF) is remaining firmly rooted in the photosphere and below, the other end is extended and stretched out by the radial expansion of the solar wind. Within 1 to 2 solar radii, this complex and highly variable field is reduced to a simple, radially directed one. The solar rotation winds up the field lines to Archimedean spirals. Thus with increasing radial distance, the original magnetic field becomes more and more toroidal. Solar rotation bends this radial pattern into an interplanetary spiral shape. Observations have revealed that the geometry of the solar magnetic field is far more complicated than that of the geomagnetic field. The Sun has a general weak dipole magnetic field with north and south magnetic polar strength of about T and multipolar magnetic field configuration during the maximum phase. The magnetometer on the pioneer 10 and 11 measured interplanetary magnetic field for radial distances between 1 to 11 AU (Smith and Barnes, 1983). On an average, the magnetic field strength can be represented as a Parker spiral (Parker, 1958). The

14 Chapter 1: THE HELIOSPHERE 20 radial distance variation of the strength of magnetic field for the ecliptic plane vary as l/r at larger radial distances. The average IMF strength is approximately 5.5 nt at 1 AU and decrease to 0.4 nt by r = 10 AU. Smith and Barnes (1983) also showed that the azimuthal component of IMF varies on average as r- 1, in conformance with the theoretical expectation. The Ulysses spacecraft, in its high latitude journey, recorded the magnetic field upto 80 0 helilatitude and found uniform magnetic field whose intensity does not change from equator to poles. The IMF, observed at 1 AU, is a weak field, varying in strength from 1 to 37 nt, with an average value of ",6 nt. Along the plane of the solar magnetic equator, the magnetic fields which act in the opposite direction run parallel to each other, separated by a thin heliospheric current sheet (RCS). The RCS is tilted due to the offset between the solar rotational and magnetic axes and warped due to the presence of a quadrupole moment in the solar magnetic field (Girish and Nayar, 1989). The IMF is a vector quantity with three directional components, two of which (Bx and By) are oriented parallel to the ecliptic. The third component (Bz) is perpendicular to the ecliptic and is created by waves and other disturbances in the solar wind. During the magnetic field reversal, usually just after the solar maximum, the solar minimum configuration of dipole component disappears and the RCS moves to higher latitudes. The inclination of the RCS is closely related to the phase of the solar cycle from low to high inclination between solar minimum to maximum. Earth crosses the current sheet at least twice during each solar rotation, sometimes more than two if the current is wavy enough. These RCS crossings are observed at earth as IMF sector boundaries, when the average polarity of the IMF changes sign. The warping in the north-south RCS is confined to relatively low latitudes during activity minimum but become accentuated with increasing activity and even reach the polar regions during activity maximum.

15 Chapter 1: THE HELIOSPHERE Coronal holes and the solar wind Coronal holes are characterized by low density cold plasma. The temperature is about half a million degrees colder that of in the bright coronal regions with unipolar magnetic fields. The magnetic field lines emerging from coronal holes are directly connected to the distant interplanetary space. In contrast, the bright coronal regions are filled with hot, higher density plasma and they are related to closed line regions. The magnetic field lines emerging from these regions return to the solar surface within a couple of solar radii. One of the major discoveries of the Skylab mission was the observation ofthe extended coronal regions in X-ray solar images. Near solar minimum, coronal holes cover about 20 % of the solar surface. The polar coronal holes are almost permanent feature, whereas the low latitude coronal holes last only for few solar rotations. The low latitude coronal holes are usually connected with the polar coronal holes. The interplay between the inward pointing gravity and outward pointing pressure gradient force result in the rapid outward expansion of the coronal plasma along the open magnetic field lines. At low latitudes, the direction of the magnetic field is far from radial. Therefore, the plasma cannot leave along magnetic field lines. At the base of the low latitude coronal holes, the magnetic field direction is not far from radial, and the expansion of the hot plasma can take place along open magnetic field lines without much resistance. This expansion greatly reduces the density of the plasma on open field lines, which appears as darker regions in the X-ray images. It should be noted that the development of the coronal holes during the solar cycle is closely related to the evolution of the solar magnetic field. The most common coronal structures seen on eclipse photographs are helmet streamers, seen as bright elongated structures, which are wider near the surface, but tamper off to a long, narrow spike at higher altitudes. The base of a helmet streamer often contains a darker cavity, in which bright prominence can sometimes be seen. Streamers appear brighter than the

16 Chapter 1: THE HELIOSPHERE 22 rest of the corona, because they are higher density regions and hence the scattering of solar radiation is enhanced. 1.4 Solar activity and solar wind The solar atmosphere is periodically disrupted by magnetic fields that stir things up, creating the active Sun. During increasing solar activity, the granulations give way form small dark areas known as pores and several pores unite together to form the sunspot. Sometimes sunspots occur in isolation, but often they arise in clusters called sunspot groups. The sunspot number reveals the solar cycle and the solar rotation. Like other solar transient features of the active Sun, the average number and location of sunspots vary in a fairly predictable cycles. On an average, the sunspot cycle lasts approximately 11 years. Maximum number of sunspots appear at the solar maximum and minimum number of sunspots appear at the solar minimum. The solar wind undergoes complicated but typical changes during the solar activity cycle which are more pronounced in the inner heliosphere. The dynamics of the solar wind is dominated by corotating solar streams and coronal transients, which are dominant over the inner heliosphere. The solar wind instruments, in the Ulysses, made a complete 3D observation of the solar wind plasma during the solar minimum phase. These observations show that the high speed polar wind that flow from the polar coronal holes spread out to lower latitudes. The polar streams possess micro streams of 3.3 day period, which may be attributed to the honey comb structure in the polar region. As we move to the solar maximum, the polar coronal holes, equatorial streamers belt separation and the assumption ofaxi-symmetry becomes less and less accurate descriptions of the global heliosphere. The streamers are no longer confined to a particular solar longitude. The ReS becomes more inclined with

17 Chapter 1: THE HELIOSPHERE 23 respect to dipole axis of the Sun as the solar activity cycle progresses, and loses its equatorial symmetry. Ulysses observations show that the slow and mixed solar wind extends to higher latitudes at solar maximum (McComas et al., 2001). A few remote measurements of some solar wind properties at high latitudes have been made closer to Sun by using radio waves scattered from fluctuations in the interplanetary medium electron density. This interplanetary scintillation (IPS) gives information on the line of sight region that is approximately 0.4 to 1 AU. Cole (1973) used this technique to show that the solar wind is slowest in the equatorial plane ( 400 kin S-I) and the fastest in the polar regions 600 km S-I. The systematic variation of 3-D structure of the solar wind according to the phase of the sunspot activity was observed by interplanetary scintillation (IPS) method (Kojima and Kakinuma, 1987). Major minimum-speed regions tend to be distributed along the neutral line through out the the solar cycle. When the solar activity is low, the Ininimum speed regions are on the wavy neutral line near solar equator. Even in the active phase, they are distributed around the neutral line which has a large warping. As the polar high speed regions expand towards the equator in the descending phase of the solar cycle, the latitudinal gradient of solar wind speed increases at the boundaries of the low speed belt and the breadth of the low speed region decrease. Similarly, when the high speed regions shrink poleward during the ascending phase of the solar cycle, the latitudinal gradient decreases and the width of the low speed region increases. When the solar activity is minimum, the polar field and the high speed region attain their maximum expanse towards equator, thereby squeezing the sector field into a narrow region and converting the neutral line into a flat structure. In the solar maximum phase, the polar high speed region contracts to a small region around the pole and disappear when the solar activity is at a maximum. The sector field expands up to high latitudes and consequently the neutral line has

18 Chapter 1: THE HELIOSPHERE 24 a large latitudinal amplitude. The sector field is not the source of the high speed wind unlike the polar field region. However, the solar wind speed from the sector field region is slightly higher than the neutral line. In addition to recurrent solar wind variations, there are many non-recurring coronal transients. 1.5 Periodic variations in solar wind The heliosphere is continuously changing due the high variability in the energy ouput of the Sun. The periodicities exhibited in the interplanetary medium arise because of the activities taking place on the surface of the Sun. These periodicities are observed by the spacecrafts due to the solar rotation (Prabhakaran Nayar, S.R., 2006). The observed periodicities in the heliosphere are grouped into short-term, mid-term and long-term types. Short-term periods are associated with the active solar longitudes or magnetic field structures (Nayar et al., 2001). The characteristics of the active regions are brought out to the interplanetary medium due to solar rotation and the outflow of solar wind. The uncertainties in the rotation period depend on the lifetime, latitudinal and longitudinal extensions of the active region. The mid-term periods, mainly 154 day, 180 day, 1 year and 1.3 year are associated with the solar dynamo, the emergence of magnetic flux and the excursion of Earth in heliolatitude (Nayar et al., 2001, 2002; Mursula et al., 2003a). The long term periods, in general, are related to the evolution of large scale solar magnetic field and observed as sunspot cycle, dipole cycle, global solar cycle or Hale cycle (Juckett, 1998). The 1.3 year period could be associated with the evolution of the differential rotation of the Sun (Paularena et al., 1995). The most prominent period near 27 days in the solar wind arise mainly because of the solar rotation. The 27 day period is most prominent in the declining phase of the solar cycle (Nayar et al., 2001). Due to the relaxation of magnetic field

19 Chapter 1: THE HELIOSPHERE 25 during the declining phase, the corona opens way for solar wind to move directly to interplanetary region with high speed. Normally, the coronal opening extend both in longitude direction as well as in the latitude direction, the opening exist for more than 15 months. The 13.5 day periodicity has its origin mainly from two high speed streams per solar rotation, detectable at near 1 AU, occur in the declining phase as well as around solar maximum. The small scale periods accounts for the coronal transients like CME, flares, corotating interaction region streams and Alfven waves. 1.6 Out of ecliptic observations The features of solar wind have been studied in detail using in situ measurements for more than three decades. All these observations, however have been confined to the near ecliptic region. The solar wind velocity has been observed by many spacecrafts starting from Mariner 2. Since this plane is inclined to the solar equator by about 7.5 0, the Earth's annual motion enables us to sample only a small range in heliographic latitude. The observations of solar wind parameters by Voyager 2, WIND and IMP 8 at different positions, phase and radial distance help us is to elucidate the solar wind evolution over a long period of time. Figure 1.2 depicts the solar wind variation at some time but at different interplanetary distances. The IMP 8 observes solar wind near to the Earth around 20 E, 'Wind' within Earth's magnetosphere and Voyager 2 in the interplanetary medium near to heliopause. The large amplitude fluctuations in the solar wind speed have been found as the predominant structure near the Earth orbit. The fluctuations of solar wind stream structure have its origin rooted to the spatial as well as temporal variations taking place in the corona. The Pioneer 10 showed a clear decrease in solar wind speed as the radial distance increased from the Sun. It is clear from the Voyager 2 observation that the velocity time profile at

20 Chapter 1: THE HELIOSPHERE Voyager - 2 ( ) (33AU-74AU) L.----'-_...L..._ I_----L.._-'--_-' l._----'-_-'-_"" '-_---'---' ] u o -(]) ;;> 800 Wind ( YOO 1AU Observation 400 IMP ear Earth Observation L-_L-_L----JL JL '_----'_-----"_-----"_----'-_----'-_---' Year Figure 1.2: Solar wind observation Wind, Imp8, Voyager2 greater radial distance from the Sun is smoother than when seen closer to the Sun. The near Earth observation brings out both coronal transients and solar cycle dependent features compared to distant observation which mainly exhibit long term variations of

21 Chapter 1: THE HELIOSPHERE 27 solar cycle effect. Definite changes in the time profile of solar wind speed is ob::;ervccl as the radial distance increases from the Sun. One major cause of this change is the interaction between the adjacent solar streams in the near Sun observation in the ecliptic position. The average solar wind speed beyond 50 AU is about 440 km s-1, a clear decrease of 30 km S Voyager observations Voyager I and Voyager II space probes, after exploring outer planets Jupiter, Saturn, Uranus and Neptune, are heading towards the outer boundary of the solar system in search of the heliopause, the region where the Sun's influence wanes and the beginning of interstellar space can be sensed. The heliopause has never been reached by any spacecraft; the Voyagers may be the first to pass through this region. The two spacecrafts should first cross an area known as the termination shock, where 11l<' solar winds slows down, which is the first indication that it is nearing the heliopause (Krimigis et al., 2003). Voyager 1 passed through the termination shock during December After the passage through termination shock, the spacecraft has enten'd the heliosheath, the region beyond the termination shock and now will be opcrntillg in the heliosheath environment. The strongest evidence that Voyager 1 has passed through the termination shock into the slower, denser solar wind beyond is that its measurement of an increase in the strength of the magnetic field carried by the solar wind and the inferred decrease in its speed. Physically, this must happen wlteiwv(~], the solar wind slows down, as it does at the termination shock. The heliosheath exploration phase ends with the passage through the heliopause which is the outer (~xt<'ll1 of the solar magnetic field and solar wind. The thickness of the heliosheath is UII(:<'1'1 <lill and can be tens of AU thick taking several years for the satellite to travers(~. Pilssag(' through the heliopause begins the interstellar exploration phas(~ with tlw SP;l(:('('}"ilfl

22 Chapter 1: THE HELIOSPHERE 28 operating in an interstellar wind dominated environment. The Voyagers should cross the heliopause 10 to 20 years after reaching the termination shock. The Voyagers have enough electrical power and thruster fuel to operate at least until By that time, Voyager 1 will be 19.9 billion kin from the Sun and Voyager 2 will be 16.9 billion km away. Eventually, the Voyagers will pass other stars Pioneer 10 and Pioneer 11 Launched on 2 March 1972, Pioneer 10 was the first spacecraft to travel through the Asteroid belt, and the first spacecraft to make direct observations and obtain close-up images of Jupiter. Famed as the most remote object ever made by man, the heliocentric radial distance of Pioneer 10 had been greater than that of any other manmade object. But late on that date Voyager l's heliocentric radial distance, in the approximate apex direction, equaled that of Pioneer 10 at AU. Thereafter, Voyager l's distance exceeds that of Pioneer 10 at the approximate rate of AU per year. The spacecraft made valuable scientific investigations in the outer regions of our solar system until the end of its mission on 31 March The Pioneer 10's weak signal continues to be tracked by the Deep Space Network (DSN) at JPL, Goldstone California, and Madrid, Spain as part of a new advanced concept study of chaos theory. Pioneer 10 is headed towards the constellation of Taurus. Pioneer 10 is in the deep space mission and it has obtained in situ measurements of the gas and dust surrounding the Sun for almost two solar cycles. Plasma analyzer measurements out to 50 AU showed that the mean velocity is constant 430 km S-I, the mean density decreases as R- 2, and that the terminal shock has not been encountered (Day 1, 1993). The magnetic field is drawn out to form Archimedean spirals in the ecliptic plane and dipole like asymmetry in the polar directions as predicted by the Parker model. Galactic cosmic ray measurement of the intensity and the radial

23 Chapter 1: THE HELIOSPHERE 29 gradient indicate a "modulation boundary" between 70 to 100 AU of the Sun. In fact all the measurements to date indicate the spacecraft is within the heliosphere. Launched on 5 April 1973, Pioneer 11 followed its sister ship to Jupiter (during 1974), made the first direct observations of Saturn (during 1979) and studied energetic particles in the outer heliosphere. The Pioneer 11 Mission ended on 30 September 1995, when the last transmission from the spacecraft was received., Its electrical power source exhausted and the spacecraft could no longer operate any of its scientific instruments, nor point its antenna toward Earth. The heliographic latitude of Pioneer 11 was greatly increased to 16 north, approximately double the maximum of excursion to which in-ecliptic spacecraft had previously restricted. Pioneer 11 was the first spacecraft to detect the spatial dependence of sector structure and found the disappearance of sector structure at high latitude Ulysses and global heliospheric structure One of the main goals of the Ulysses mission was to determine the global structure of the heliosphere, in particular with regard to the distribution of solar wind plasma, and its frozen-in interplanetary magnetic field. The three-dimensional structure of the heliosphere is characterized by a basic north-south symmetry, and is dominated by the presence of the fast solar wind from polar and high-latitude regions that expands to occupy a large fraction of the heliospheric volume. Observations made by Ulysses during its rapid pole-to-pole transit near the perihelion have revealed that the transition from slow to fast wind is surprisingly abrupt. The Ulysses magneticfield measurements indicate that the radial field component, and consequently the magnetic flux, are independent of heliolatitude. This has led to the unanticipated conclusion that the magnetic-field pressure controls the solar-wind flow near the Sun, driving a non-radial expansion by more than a factor of 5 from the polar coronal holes.

24 Chapter 1: THE HELIOSPHERE 30 Another characteristic property of the fast high-latitude wind that was revealed by Ulysses is the continual presence of large-amplitude transverse waves in the magnetic field. The wave amplitudes typically equal or exceed the magnitude of the background field and, as such, are an example of strong turbulence. Although the spiral magneticfield structure predicted by earliest models of the solar wind is preserved to the highest latitudes, a significant departure from the Parker model has been discovered. Ulysses observations of the solar radio bursts, which are relatively rare at solar minimum, provide a capability for remote-sensing of large-scale magnetic-field structures. Ulysses discovered that the recurrent effects in energetic-particle and cosmic-ray intensity extend to much higher latitudes than the CIRs themselves. Under the influence of the magnetic configuration of the Sun, the angle between the solar rotation and magnetic axes changes with the solar cycle. Near solar minimum, the magnetic and rotation axes are nearly aligned, so that the CIRs are restricted to relatively low heliographic latitudes. The details of the scientific instruments housed in the Ulysses and the mission objectives are described in chapter Indirect observations Comet observations The plasma tails of the comets act as a wind socks, whose orientation can be used to infer the solar wind velocity throughout the heliosphere. The solar wind magnetic flux and magnetic structure can be collected from the extent of the Comet's hydrogen cloud as observed in its Lyman alpha radiations, and perhaps also from the occurrence of disconnections events which the entire tail becomes detached from the comet's head. The sensitivity of charge exchange processes to properties of both the solar wind and the comet leads to many observable effects. The observation of helium emission from

25 Chapter 1: THE HELIOSPHERE 31 comets can therefore be used to probe the characteristics of the comet solar wind interaction. Emission cross sections are very sensitive to velocity effects. Every charge exchange reaction has its own particular behavior with respect to the velocity, so that line emission ratios - in the case the ratio between the Hen 30.4 nm and the HeI 58.4 nm line - can be used for velocimetry. Absolute intensities can be used to determine local heavy ion fluxes. Other information on heliospheric structure both in and out of ecliptic plane has been gathered from analysis of variations in the cosmic ray flux and from radar probing, and phase delays in radio signals transmitted from interplanetary spacecraft. Most recently the study of long kilometer plasma radio signals from the interplanetary disturbances themselves is providing important information on how such transients propagate from the Sun to Earth's orbit and beyond Interplanetary scintillation technique The interplanetary scintillation (IPS) observations can be used for the determination of the structure and dynamics of the solar wind (Manoharan, 1993). The most important information of solar wind behavior at high latitudes has been obtained from interferometric studies of radio wave propagation from cosmic radio sources (Manoharan et at., 2000; Joshi et at., 2006). In early 1950's, an apparent broadening in the radio sources as found when viewed in the vicinity of the Sun, the effect could be explained by irregularities in the outer coronal structures (Hewish et at., 1964). Observations of small diameter cosmic radio sources also exhibit flux variations, or scintillations, on time scales of Hz attributed to interference between the scattering patterns caused by the plasma density variations. Movement of the irregularities with the solar wind leads to drift of the interference pattern across an observer on earth, causing the flux variations. Spatial coherence of the pattern can be verified by comparison of scintillation observed at antennas by several tens of kilometers. Solar

26 Chapter 1: THE HELIOSPHERE 32 wind velocity measured using IPS is compared with in situ spacecraft measurement shows good agreement, provided the solar wind flow is dominated by steady velocity structures rather than by relatively short-lived transients. Mapping of the solar wind with IPS provides the solar wind structure in terms of latitude and longitude without much ambiguities. Sime and Rickett (1978) first introduced the two dimensional speed map, and these maps have been used for studying the relationships between the solar wind speed and coronal parameters such as coronal brightness and magnetic field distribution(rickett and Coles, 1982). The variations seen in these graphs provides the general trends in the large scale structure. 1.8 Conclusion This chapter provides introductory material for the following chapters. Proof of the existence of the solar wind was one of the first great triumphs of space age, and enormous details has been brought out about the physical nature of the solar wind plasma in the last 45 years. Nevertheless, our understanding of the solar wind is far from complete. We are yet to know the heating, and the solar wind acceleration mechanisms. We are still confused about the origin of slow solar wind, whether it is from quasi-stationary processes or from small coronal transients. On smallest scales, the properties of the solar wind are found to be randomly varying and turbulent. At intermediate level, the dominant features show sharp variations when encountered along various boundaries. On largest scales, there is surprising amount structure, with gross variations in all directions. The solar system is bathed in the hot gale wind that blows from the Sun, carves a cavity in the interstellar space called a heliosphere. Various spacecraft observations have helped to find the shape and to estimate the contents of the heliosphere.