CHAPTER I

HELIOSPHERIC CURRENT SHEET Al~D

INTERPLANETARY l4EDIUM

page

1.1. Introduction 1

1.2. IJ.\1F sector structure 2

1.3. Solar wind at 1 AU 8

1.4. Large scale features of the solar corona 19

1.5. Heliospheric current sheet 28

I,• i

1

CHAPTER I

HELIOSPHERIC CURRENT SHEET AND INTERPLANETARY MEDIUM

1.1. Introduction

The expanding solar corona carries the plasma and

magnetic field of the solar atmosphere to great distances

out in the solar system and provides the interplanetary

plasma and magnetic field (IMF). The solar wind and IMP have

extensively been studied during the past three decades. The

interplanetary medium shows a large variety of short time

scale and long time scale fluctuations depending on the

phase of the solar activity cycle. The IMF is highly

fluctuating on a short time scale, but on an average of few

hours or more, it is ordered into large scale features.

Sector structure is one such large scale feature, with

average magnetic field directed either away from the sun

(away sector) or toward the sun (toward sector) roughly

along the Archimedean spiral direction defined by the zero­

order model of Parker (1958, 1963). The sector str~cture

observed in the interplanetary medium (Ness and Wilcox,

1964; Wilcox and Ness, 1965) is now understood as a result

of the presence of a heliospheric current sheet separating

the heliosphere into two hemispheres of opposite dominant

2

magnetic polarity. The heliospheric current sheet is

observed near earth as IMF sector boundary at which the I~1F

field-direction reverses within few minutes to few hours.

The number of sector boundaries observed near earth during a

solar rotation period depends on the warpings of the HCS

which in turn is determined by the complexity of the solar

magnetic field. Thus the study of the morphology and

periodic variations of the sector structure can give an idea

of the complexity of the solar magnetic field. The study of

the solar activity and its signature on the interplanetary

medium is important in understanding the solar terrestrial

relationship. In recent years, a wide variety of techniques

and methods have been used to understand the solar processes

and their relationship \vith the interplanetary medium. The

aim of this chapter is to introduce some of the basic

concepts available in literature on IMF sector structure,

solar wind at 1 AU, large scale features of solar corona and

heliospheric current sheet, so as to give a background to

the problems discussed in the following chapters.

1.2. I~W sector structure

The measurements of IMF polarity have been carried

out using the magnetic field data obtained from earth

orbiting satellites since 1964 (VJilcox, 1968; Couzens and

3

King, 1986). An elt8TnAtp ~e~hod of infprring the onl@rity

of IMF near earth was discovered by Sva1gaard (1968) and

Mansurov (1969) making use of the polar cap geomagnetic

field observations. From this method, one can infer the I~lF

polarities back to 1926 (Svalgaard, 1972; Mansurov and

Manurova, 1970; Sva1gaard, 1976; Matsushita and Trotter,

1980). The inferred IMF polarity using ground based

magnetograms of polar stations [Thule (76 0 33'N) and Vostok

(84 0 54'S)J show a high degree of agreement with those

determined using satellite observations and hence used for

solar-terrestrial studies (Russel et al., 1975; Wilcox et

al., 1975). Mori and Nagashima (1979) developed a method to

infer IMP sector polarity from the cosmic ray north-south

asymmetry, utilising ground level cosmic ray observations.

Nefdner (1982) describes a method to infer IMF sector

boundary passage from aa indices of geomagnetic activity

after studying 72 disconnection events of 29 comets

extending back to 1892. The mean field observations of solar

magnetic field (magnetic field of sun as a star) is also

another means to infer the IMF polarity near earth (Severny

et ~l., 1970; Scherrer et al., 1977a,b).

The IMF sector structure and sector boundary is

related to a large number of solar-terrestrial phenomena.

The IMF sector structure re~lect the large scale solar

magnetic field (Svalgaard and Wilcox, 1975, 1976a,b; Wilcox

and Scherrer, 1981). So one can infer the large scale solar

magnetic field by studying the properties of sector

structure. The interplanetary plasma parameters like

magnetic field intensity, particle density, temperature,

solar wind velocity, cosmic ray flux, interplanetary ion

flux and solar activity parameters like solar flare

occurrence, calcium plage occurrence, green corona and K­

corona brightness, solar radio emission etc. are found to be

organised around the IMF sector boundary (or RCS) in

charactEristic form (Wilcox, 1968, 1977). Apart from this,

several terrestrial phenomena are also found to be organised

around the IMF sector boundaries (Taylor, 1986; Svalgaard,

1977; Arora and Rangarajan, 1981).

Heliolatitudinal variations in I~~

Most of the measurements of INF is linited to

heliographic latitudes +7.25 0 related to annual motion of

earth around the Sun. There are also few out of ecliptic

spacecraft measurements of INF. Rosenberg and Coleman (1969)

proposed a dominant polarity effect of I~lF from analysis of

IMF measurements from different space-crafts between 1962

and 1969. In their model, the dominant polarity of IMF in

the northern or southern heliohemisphere coincides with the

polarity of solar polar magnetic ~ield of the same solar

5

hemisphere (Rosenberg, 1970). This model also predicts a

reversal of the dominant polarity of I~lF ln a particular

heliohemisphere with the reversal of the solar polar field

during the post sunspot maxima periods and the effect of 22

year heliomagnetic cycle on the IMF sector structure (Wilcox

and Scherrer, 1972; Hedgecock, 1975; Smith et al., 1986).

Latitude dependence of IMF polarity during 1975-1981 is

studied by Villan~e and Mariani (1983) using satellite

observations. Eventhough, the dominant polarity effect

proposed by Rosenberg and Coleman (1969) is apparent in near

earth observations for few years, in general the statistical

support is not satisfactory when long term IMF observations

are analysed (~loussas and Tritakis, 1982; Tritakis, 1984a;

Xanthakis et al.1981; Neidner,1982). Behannon et ale (1989 )

used source surface magnetic field data computed from

potential field modelling of photospheric magnetic field

observations for the period 1977-1984 to show the predomi-

nant polarity variations of the IMF with heliographic

latitude within +70 0. Pioneer and Voyager missions provide

out of ecliptic observations of IMF up to 27 0 north of

heliographic equator and up to a radial distance of 25 AU

from the Sun (Behannon et al., 1989; Mihalov et al., 1989;

Smith et al., 1978).

6

Norphology of INF sectors

Generally two or four IMF sectors are observed per

solar rotation near 1 AU and occurrences of six sectors are

rare. The complexity of the sector structure vary with solar

sunspot cycle. T\vO sectors are more frequently observed in

the high solar activity periods (Svalgaard et al., 1975;

Sawyer, 1974) and four sectors are observed near minimum and

post sunspot maxima periods. A typical large scale structure

of IMF have life times up to a period greater than one year

(Svalgaard et al., 1975). Behannon et ale (1989) obtained a

mean life-time of 6 to 7 solar rotations for the

heliospheric sector structure. The complexity of the sector

structure vary with heliolatitude of observations (Sawyer,

1974; Behannon et al., 1989). The sector boundary also

prefers to occur in certain heliolongitudes or certain days

in Bartel's rotation period (Tritakis, 1979a, 1986). The

number of sectors observed per solar rotation also depend on

the relative contribution of various magnetic multipoles

present in the solar magnetic field (Schultz, 1973; Levine,

1980a; Hoeksema, 1984). Wang Shui and Fang Lizhi (1979)

discussed the IMF sector structure in terms of large scale

spiral waves in the solar equatorial plane. Plyusnina (1985)

observed a north-south asymmetry in the correlation between

IMF and photospheric magnetic field structures in the

7

opposite heliohemispheres. Luhmann et ale (1987) reports on

the asymmetry in the IMF strength about the heliographic

equator from spacecraft observations in the interplanetary

medium.

Periodicities in the Il4F polarity

Observations of IMF structure near lAU generally

show a 27 day recurrence period. Sometimes during solar

maxima) a 28 day recurring IMF structure is also found to

coexist with the 27 day pattern. The recurrent pattern of

the IMF change with phase of the sunspot cycle and it

remains relatively stable with a period. of 27 days during

the declining phase of the solar cycle (Svalgaard 1972;

Svalgaard et al., 1975). Hoeksema and Scherrer (1984)-- also

reports a periodicity of 27 days and 28 days associated with

equatorial dipole component of the coronal magnetic field.

Sheeley and Devore (1986) and Nash et ale (1988) proposed a

model to explain the 27-28 day periods in IMF sector

pattern. A comparison of coronal and IMF recurrence period

during 1964-1982 was reported by Parker (1987). Using

inferred IMP polarities, between 1926 and 1986 Gonzalez and

Gonzalez (1987) obtained various periodicities about 27.5,

13.5, 9.1 and 6.8 days associated with the evolution of IMF

sector pattern. They also found periodicities of about 3.7

8

and 1.5 years associated with change ln the dominant

structure of IMF from a well defined 2 sector to 4 sector

structure. Xanthakis et ale (1981) also reports on the power

spectral analysis of inferred IMF polarity (1926-1981).

Kotov and Levitskii (1984) found periodicities in the range

26-30 days in the evolution of mean magnetic field of the

Sun and IMF polarity from a power spectral analysis of solar

and interplanetary data. Similar periodicities are also

found to exist in geomagnetic and sunspot data (Gonzalez and

Gonzalez, 1987; Villante and Francia, 1986, 1988).

1.3. Solar wind at lAU

rrhe concept of coronal expansion out into

interplanetary medium emerged as early as 1900 and gained

support from the observations of comet tails (Vlilcox, 1976;

Birkeland, 1909; Biermann, 1951). The spadecraft observa-

tions in early sixties confirmed the constant solar wind

flow in the near earth environment. Basically there are two

types of solar wind flows observed near lAD associated with

distinct solar origin and properties viz., low speed and

high speed solar wind. This background flow of solar wind is

occasionally modified by the transient flows from features

in the Sun like solar flares or coronal mass ejections.

Table 1.1 shows typical characteristics of three types of

9

Table 1.1

Comparison of solar wind parameters in three

types of flow observed at 1 AU

Parameter High speedsolar ".,rind

Lo\-v speedsolar wind

Solar wind relatedtosolar transients

km s -1 700 380 440v,

-3 4 15 10n, em

8 -2-1 3 5 5nv, 10 em s

2 ,dyne -2 3 3 3nmv em

nv(mv 2/2+rnHG/r),

-2 -1erg em s

T 105 Kp'

T 105 Ke'

n v/n vp p

2

2.3

1.0

.05

2

0.7

1.3

.02

2

0.6

0.5

.10

Ionisation temperatures

6TO' 10 K ?

6TFe , 10 K ?

-3 -2 -1 3Q 10 erg em se'

B, nT 6

Field topology open

= 8 kT /B2 1p

B /nv, 10- 3nTem3skm-1 1r

2.1

1.6

<3

3

?

2

0.3

up to 3.4

up to 17

1

9

closed

0.3

2

10

solar wind flows (Neugebauer, 1983).

Low speed solar wind

Low velocity solar wind with speed 300-400 kms- l

are generally associated with the closed magnetic field

regions in the Sun with enhanced electron density viz.,

coronal streamers or heliospheric current sheet (Gosling et

al., 1981; Esleev and Filippov, 1988). Solar wind properties

show substantial variability, in most of the parameters in a

low speed solar wind flow (Feldman et ~~., 1977). Proton

temperature and momonentum flux density observed near lAU in

the low speed solar wind show an inverse relationship (Lopez

and Freeman, 1986; Sastri, 1987). The helium abundance also

observed to be lower in the low speed solar wind (Borrini et

al., 1981; Feldman et al., 1981; Gosling et al., 1981). The

slow speed solar wind carries away the positive angular

momentum from the Sun and it is released from the Sun at

significantly lower Alfven radius than the fast solar wind

(Schwenn, 1987; pizzo et al., 1983). Kojima and Kakinuma

(1987) reported the solar cycle variation in the loV! speed

solar wind with respect to HCS using IPS (Interplanetary

Scintillations) observations. Multi-spacecraft study of the

spatial structure of solar wind velocity during 1985-87 is

recently reported (r1iyake et al., 1988). Lyubimov and

11

Pereslegina (1985) discuss on the chromospheric and coronal

sources of low speed solar wind.

High speed solar wind

Examination of solar wind data accumulated by

different spacecrafts reveal continuous presence of various

types of structures in the expanding solar wind. One of theQ

is the high speed stream which was identified in the Mariner

2 data of 1962 (Neugebauer and Snyder, 1966). A high speed

plasma streaQ is characterised by a large increase in the

solar wind velocity lasting for several days. The flow tends

to come from slightly east of the Sun as the speed begins to

rise and form a westerly as the solar wind speed approaches

its peak (Ness et al., 1971; Siscoe et al., 1969). Different

definitions have been used to identify and catalogue high

speed streams (Intrigilator, 1973, 1977; Bame et al., 1976;

Gosling et al., 1976; Broussard et al., 1977; Lindblad and

Lundstedt, 1982, 1983; Mavromichalaki, 1988). Generally, two

different classes of high speed streams observed near 1 AU

are (1) cororating streams flowing generally from coronal

holes and (2) transient streams associated with solar flares

(Iucci et al., 1979; Tsurutani et al., 1987). The basic

features of cororating high speed streams are:

12

(a) The proton density (n) rises to unusually high values

near the leading edge of the streams persisting for

one day and generally having an inverse relationship to

bulk speed (V).

(b) The IHF magnitude (B) is proportional to bulk speed

with constant IMF polarity throughout the stream except

for fluctuations lasting for few hours.

(c) Proton temperature (Tp) pattern is similar to bulk

speed.

The properties of flare-related stream are:

(a) All the interplanetary parameters show simultaneous

increase, possibly denoting radially outco~ing fast

shocks.

(b) V, Tp and B show large fluctuations in the maximum

solar wind region. The polarity of the magnetic field

often show inversions.

(c) Tp behaviour tends to depart from solar wind behaviour.

The high speed stream - IMF sector structure

relation is investigated by many authors (Gosling et al.,

1976, 1978; Hundhausen, 1977; Sawyer, 1976; Sheeley and

Harvey, 1981; Sheeley et al., 1977; Obridko and Shelting,

1987a,b). The high speed steams of solar wind are generally

found to be associated with single IMF polarity. There are

13

IMF sectors without high speed streams and sometimes more

than one high speed stream are found within a IMF sector.

The properties of high speed streams are found to evolve

with solar cycle (Intrigilator, 1977; Bame et al., 1977).

The preferred Bartel's day distribution of the occurrences

of high speed streams during cycle 20 and 21 has been

investigated (Lindblad, 1981; Rangarajan and Mavromicha1aki,

1989). Streams observed during the ascending phase and

maxima of the sunspot cycle are generally short lived and

non-recurring while stable recurrent high speed streams are

observed during the declining period of sunspot cycle. The

duration of high speed solar wind stream can vary from few

days to nearly two weeks as observed near 1 AU. The

amplitude of the high speed streams vary between 450-800

-1kms . Broader streams are observed during low solar

activity periods. Both equatorial coronal holes and

equatorward extension of polar coronal holes are identified

to be the source of high speed solar wind streams observed

near the earth (Hundhausen, 1977). The amplitude or velocity

maxima of high speed stearns are observed to be correlated

with size, area and geometrical divergence of magnetic field

lines of coronal holes (Nolte et al., 1976; Schwenn et al.,

1978; Eslevich and Fi1ippov, 1986). The large scale magnetic

field of the solar corona is observed to control the three

dimensional density structure of the corona and associated

14

flow of solar wind into the interplanetary medium (Levine,

1978, 1980b; Levine et al., 1977; Pneumann, 1976; Pneumann

et al., 1978; Burlaga~! al., 1978; Hundhausen, 1977).

Shukhova et al. (1987) discusses on the open loops of IMF

observed in the ecliptic high speed solar wind. Burlaga

(1986) reviews our current knowledge on the structure and

dynamics of corotating and transient solar wind streams in

three dimensions. Henning et al. (1985) made a study on the

observation of solar flares and related high speed streams

near 1 AU in relation with the structure of the heliospheric

current sheet.

Apart from the observations of the solar wind near

sun's equatorial plane by earth orbiting satellites IPS

(Interplanetary scintillations) observations provide

information regarding the solar wind flow at different

heliolatitudes with in +70 0 of the solar equator (Sime and

Rickett, 1978, 1981; Coles et al., 1980; Rickett and Coles,

1983; Kojima and Kakinuma, 1987). Coles et al. (1980) found

from IPS observations that the solar wind flow from the

polar regions of the sun vary with sunspot cycle. The flow

properties of the solar wind from the equatorial and polar

region of the sun differ in many aspects) ,contributing to

the observed solar wind velocity variations near 1 AU ~vithin

ecliptic (Simon and Legrand, 1987). It is observed that the

K-corona brightness or electron density distribution in the

15

solar corona is inversely correlated with solar wind speed

(Sime and Rickett, 1978; Rickett and Coles, 1983). The

variation of the properties of polar solar wind with sunspot

activity is related to the corresponding change in the

poloidal solar magnetic field and the associated change in

flux tube divergence of magnetic field lines originating

from polar coronal holes (Simon and Legrand, 1986; Coles et.

al., 1980). Long term variations in the ecliptic solar wind

can be inferred from geomagnetic or sunspot data (Legrand

and Simon, 1981, 1985; Simon and Legrand, 1986, 1987;

Russel, 1975; Sargent, 1986; Silverman, 1986).

Despite considerable solar wind and related

observations a~e accumulated over the past three decades the

basic physics of the heating and acceleration of the solar

wind is yet to be understood properly. Excellent reviews on

this subject exists (Hundhausen, 1972; Hollweg, 1978;

Withrobe, 1986; Holzer, 1979; Leer et al., 1982; Leer and

Holzer, 1985; Pneumann, 1986; Marsch, 1985; Schwenn, 1987;

Leer, 1987). Apart from physical mechanisms like purely

thermal acceleration with and without extended heating,

acceleration due to Alfven wave pressure, diamagnetic

acceleration, the role of sporadic events such as spicules,

macrospicules, X-ray bright points, ephemeral magnetic field

regions and out flows seen in the EUV related to explosive

events in the Sun are all explored in this connection

16

(Pneumann, 1986; Yang and Schunk 1989).

Heliornagnetic latitudinal organisation of· solar wind speed

The concept that the three dimensional coronal

magnetic field geometry controls the spatial structure of

the solar wind in the heliosphere gained support through

several studies after the skylab period (Levine et al.,

1977; Hundhausen, 1977, 1978; Svalgaard and Wilcox, 1978;

Schultz et al., 1978). Subsequently Hakamada and Akasofu

(1981) demonstrated the feasibility of explaining most of

the solar wind speed variations (daily, 27-day, semi-annual

etc.) observed near 1 AU by assuming a positive gradient in

solar wind speed with angular distance from the heliospheric

current sheet (heliomagnetic latitude) and change in

heliomagnetic latitude of the observing point (e.g. earth or

space craft).

Zhao and Hundhausen (1981) found a relation

V(km/s) = 400 + 1000 sin2~ (1)

during 1974 between solar wind speed V and the angular

displacement ~ from a flat heliomagnetic equator inclined

approximately 30+10 deg with respect to the equatorial plane

of the Sun. Zhao and Hundhausen (1983) found a new

relationship

V = 350 + 800 * sin 2p ( 2 )

17

for 1976 where p is the angular displacement from HCS

inferred from K-corona observations and a plateau of 600

kms -1 for f3 > 135/day.

Hakamada and Munakata (1984) obtained a relation

V = 408 + 473 sin2y (3)

for the period 1976-77 wherey is the heliomagnetic latitude

of the observing point derived from HCS inferences using

potential field modelling of solar magnetic field. Coles et

ale (1980) from a study of solar wind observations using IPS

observations suggested that the helioBagnetic latitude

dependence of solar wind speed can change with the sunspot

cycle. This is confirmed by the study of Newkirk and Fisk

(1985) concerning the heliomagnetic latitude dependence of

solar wind speed using HCS inferenc8s from synoptic X-corona

data for the period 1964-1982 and subsequently by Fry and

Akasofu (1987) investigating the sa~e using source surface

magnetic field observations for the period 1976-1982. It is

found that heliomagnetic latitudinal gradient of solar wind

speed is steeper during sunspot maxir;l.Um. Bruno et ale 1986)

investigated the latitudinal gradients of different solar

wind parameters during 1976-1977 using in situ observations.

Recently Kotova et ale (1987) reported on heliomagnetic

latitude dependence of solar wind speed using PROGN07.-9

measurements.

One can tind evidence for correlati n ll between

18

solar wind speed and magnetic field strength in the source

surface in various studies (Steinolfsen, 1982; Pneumann and

Kopp, 1971; Yeh and Pneumann, 1977; Suess et al., 1977;

Munro and Jackson, 1977; Hoeksema, 1984). Suess et ale

(1984) found that the source surface field strength and

solar wind speed near 1 AU correlated better than the

heliomagnetic latitude during 1979. They also show the

importance of removing effects due to transient related

events when we study the heliomagnetic latitude dependence

of solar wind speed near 1 AU. During 1974, Hakamada (1987)

found that angular distance from the HCS is a better

organiser of the solar wind speed than the source surface

field strength. Kojima and Kakinuma (1987) investigated the

solar cycle variations in the solar wind speed in relation

with change in Hes structure using IPS observations and also

suggested the limitations in expressing solar wind speed as

a simple function~heliomagneticlatitude. Bruno et ale

(1986) and Fry and Akasofu (1987) independently showed that

solar wind speed distribution is asymmetric about the HCS. A

22 year cycle variation in the solar magnetic structure and

solar wind is also proposed (Chirkov and Sansonov, 1984).

According to Simon and Legrand (1987) the helioDagnetic

latitude dependence of solar wind change from one solar

cycle to another following the change in the maximum of

sunspot activity.

19

1.4. Large scale features of the solar corona

The white light corona becomes visible during

total solar eclipses and there is an observational record of

white light corona during total solar eclipses for nearly

100 years. Outside the eclipses it can be observed at 1-10

Re by balloon borne or satellite borne coronographs (e.g.,

080-7, 8kylab) or ground based K-corona Qeters (e.g., Mauna

Loa observations in Hawaii Islands).

The K-corona or electron corona is produced by

Thompson scattering of photospheric radiation by the

electrons of the highly ionized coronal plasma. The K-corona

is very inhomogeno~s containing a number of characteristic

structures such as streamers, arches, plumes, fine rays and

coronal holes. K-corona is observed by the Mauna Loa K­

corona meters since 1965 (Hundhausen et al., 1981). The data

is presented in the form of synoptic charts, derived from

successive daily scans in position angle around the solar

limb. Each map is a cylindrical projection of coronal

polarisation brightness (pB) plotted against solar latitude

and longitude. The pB provides a measure of integrated

electron density of the corona at a given height (1-4 R0)

above the solar surface (Perry and Altschuler, 1973;

Hundhausen, 1977). The K-corona pB map for the carrington

rotation 1614 is shown in figure 1.1 (8ime, 1988).

20

-

HIGH f,J..,TIT\££ ~ATCRi, HA~ LOA, ~11

CONiQRS (1:' K-C~ IPS) AT 1.500 R

ROT;HICN 1614, E:AST L1t13

OCGIt'NIN::; DAy 107.4, \974""1.t­o

.I

I

vo90

.. ~ -"t' ;" ......... -..~ ..

I

~r"--._...--- ---.a--I.-L-"""'-'+__--¥_-I-1I_+-+-~---=:>.{-\,,..+---I....-- +'I-

o

..

carrca..R:i ~ " 2, 3, 4, 6, 8, 10, 12 ffi X 10-8

Fig.l.l. K-corona brightness synoptic map for the carrington

rotation 1614 (Sime, 1988)

21

Due to its high temperature of the order of

million degrees, the solar corona shows line radiations from

highly ionized atoms in emission against continuous

spectrum. The most prominent lines are red (Fe X A 6374 "A)

o 0

the green (Fe XIV ~ 5303 A) and the yellow (Ca XV ~ 5694 A)

(Waldmeir, 1971). Outside the total solar eclipses these

emission lines are observed employing a Lyot coronograph

with appropriate narrow band filters of spectrograph. The

ogreen corona (A 5303 A) has observational record since 1939

(Waldmeir, 1981; Leftus and Sykora, 1982). In addition to

this, corona can also be observed in X-ray and XUV wave

lengths using mainly space crafts (Broussard et al., 1978).

oThe absorption lines of Heliulu (He I 5876 A and He I 10830

oA) and radio observations can also give inforQation

regarding the solar corona (Harvey and Sheeley, 1979).

Coronal streamers

Coronal stream.eV$ are approximately radial

structures of high electron density found between 0.5 to 10

R®. Active region steamers forDing above young active

regions are short lined, while helmet streamers lying above

quiescent prominences or extended bipolar regions live for

many months (Koutchmy, 1977). Streamers are generally

associated with magnetic field structure that is closed and

inhibit outflow of plasma (Pneuma~n et al., 1978). The

22

heliospheric current sheet which separates regions of

opposite dominant magnetic polarity in the heliosphere is

identified on the centre of the bright band of coronal

streamers observable in K-corona (Hundhausen, 1977; Korzhov,

1977). Fisher and Sime have co~ducted a study on the coronal

streamers using K-corona observations for the interval 1965-

83 (quoted by Sime, 1986) using Mauna Loa K-corona

observations. Number of streamers observed per solar

rotation change with sunspot activity with maximum number

being observed in sunspot maxinum and minimum number during

sunspot minimum. Most of the structures occur near solar

equator. The latitude of the highest streamer is In close

correspondence with the position of HCS in comparison with

other features like latitude of the highest filaments during

a solar rotation period.

The low speed solar wind flow (300-400 km -1s )

near 1 AU is associated with coronal streamers (Borrini et

al., 1981; Feldman et al., 1981; Gosling et al.,

Sastry, 1987).

Coronal holes

1981;

Coronal holes are regions of abnormally low

coronal density associated with magnetic field lines that

are 'open' to interplanetary space. They were known to exist

24

Large coronal holes are present at the solar poles

near declining phase and sunspot minimum and they shrunk in

size during ascending phase of the solar cycle and disappear

during the solar maxima. The area of the polar coronal holes

waxes and and wanes with net amount of magnetic flux present

at the solar poles (Broussard et al., 1978; Hundhausen e~· a~.)

1981; Webb and Davis, 1984). Coronal holes which occur at

equatorial or middle solar latitudes are frequently

same

regions

the

holes tend to occur in magnetic

polarity as the polar caps insamethehaving

connected to polar coronal holes, whenever sufficien~ uni­

polar regions develop at these sites. Their origin and

evolution are more difficult to predict, but they are more

stable during the declining phase of the sunspot cycle.

Coronal holes are located within large scale magnetic areas

dominated by one magnetic polarity and having a diverging

(open) field geometry (Hundhausen, 1977). They for~ over

regions where the magnetic field strength is locally high

(0.7 G to 12 G) as shown in different studies (Bohlin and

Sheeley, 1978; Harvey et al., 1982). A strong, large scale

magnetic field will usually develop before a coronal hole

and remain long after the coronal hole disappearance which

suggest that coronal hole is an evolutionary step in the

development of large scale magnetic regions in the Sun

(Hoeksema, 1984).

Coronal

24

Large coronal holes are present at the solar poles

near declining phase and sunspot minimum and they shrunk in

size during ascending phase of the solar cycle and disappear

during the solar maxima. The area of the polar coronal holes

waxes and and wanes with net amount of magnetic flux present

at the solar poles (Broussard et al., 1978; Hundhausen e/::. a.~')

1981; Webb and Davis, 1984). Coronal holes which occur at

equatorial or middle solar latitudes are frequently

connected to polar coronal holes, whenever sufficien~ uni­

polar regions develop at these sites. Their origin and

evolution are more difficult to predict, but they are more

stable during the declining phase of the sunspot cycle.

Coronal holes are located within large scale magnetic areas

dominated by one magnetic polarity and having a diverging

(open) field geometry (Hundhausen, 1977). They for~ over

regions where the magnetic field strength is locally high

(0.7 G to 12 G) as shown in different studies (Bohlin and

Sheeley, 1978; Harvey et al., 1982). A strong, large scale

magnetic field will usually develop before a coronal hole

and remain long after the coronal hole disappearance which

suggest that coronal hole is an evolutionary step ln the

development of large scale magnetic regions ln the Sun

(Hoeksema, 1984).

Coronal holes tend to occur in magnetic regions

having the same polarity as the polar caps in the same

25

hemisphere but exceptions are observed (Harvey and Sheeley,

1979). Coronal holes exist directly adjacent to disk

activity (Bohlin and Sheeley, 1978). The life time of

coronal holes has a mean value of 6 solar rotations and

range froQ 3 to 20 solar rotations (Bohlin, 1977). The long

term growth in size and decay of coronal holes occur at a

rate of about 1.5 xl0 4 km2s-1 which is consistent with the

diffusion rate of magnetic field (Timothy et al., 1975;

Bohlin, 1977). The area of coronal holes change mainly by

sporadic, large scale shift of the boundaries (Nolte et al.,

1978a,b). The boundaries of established coronal holes can be

altered by near prominence eruptions and short lived coronal

holes can be produced by these prominence eruptions and

flares (vlebb et al., 1978). The coronal holes generally

rotate rigidly with a synodic period of 27.2 days but

occasionally some differential rotation is also observed

(Timothy et al., 1975; Wagner, 1975; Bohlin, 1977; SheIke

and Pande, 1985). The solar cycle evolution of coronal holes

is explained in terms of locally flux iQbalance model of

Broussard et ale (1978). There is a systenatic heliolatitu-

dinal variation of occurrence of coronal holes during a

solar cycle but indications of sone solar longitudinal

organisation of coronal holes is also there (Broussard eteru.I')

al., 1978; Hundhausen) 1981; Svalgaard and Duvall, 1977).

The inter-relationship between coronal holes, solar wind

26

streams, IMF sector structure and geomagnetic activity

during solar cycle 20 and early half of solar cycle 21 is

given in a pictorial format by Sheeley et al. (1977) and

Sheeley and Harvey (1981).

Spatial and temporal changes in the coronal structure

The overall shape of the corona is observed to

change with the solar cycle. The ellipcity coefficient '~'

characterising the shape of the isophotes of the white light

corona is found to be small (~-0.05) during solar maximum

and large (£-0.25) during solar minima (Koutchmy, 1977). The

variation in the apparent shape of the corona arises partly

from the changes in the distribution of bright features

(steamers) and dark regions (coronal holes) observed around

the limb in white light. Streamers are distributed fairly

uniform around the limb of the Sun giving almost a circular

shape of the corona during sunspot maxima. As the solar

cycle progresses streamers tend to occur close to the solar

equator and coronal holes establishes themselves at poles.

During the declining phase and minimum of the sunspot cycle

polar coronal holes dominate the high and mid latitudes of

the Sun, while streamers are found only in low solar

latitudes. During ascending phase streaDers gradually spread

27

to higher solar latitudes and polar coronal holes \lill

shrink and eventually disappear near solar maximaUiundhausen

et al., 1981).

In the inner corona the magnetic field dominates

the plasma \lhile in the outer corona radial flow of solar

wind dominates. The coronal structure and its change

reflects that of the solar magnetic field. At miniQa the

large scale magnetic field is predoDinantly dipolar and the

open polar field lines provide large coronal holes at the

poles. Long streamers extend far out along the equator

marking the extension of surface fields by the solar \lind.

At sunspot maximum the Qulti-polar components in the

heliomagnetic fields are quite dominant and we see a bright

corona all around the limb. Active regions provide coronal

condensations and their magnetic remnants give rise to

helmet streamers. Balough (1986) and f1c Queen (1986) reviews

our current understanding on heliospheric and coronal

magnetic fields respectively.

The integrated brightness of the corona varies

with sunspot activity as known from the study of Fisher and

Bime (1984) using K-corona observations. In addition to the

observed pole-equator asymmetry in the coronal brightness

distribution, a north-south asymmetry in the same about the

solar equator is also observed near sunspot minimum as

evident from K-corona and green corona observations (Houssas

28

et al., 1982; Hundhausen et al., 1981). Tritakis et al.

(1988) investigated the long term variations in the

latitudinal and longitudinal asyn~etry of the solar coronal

structure using green corona observations during 1944-1974.

Waldmeir (1981) studied the long term variations in area of

the polar coronal holes bet~leen 1939-1980 using green corona

observations.

1.5. HELIOSPHERIC CURRENT SHEET

Heliospheric Current Sheet (HCS) is a neutral

sheet separating solar \vind flows carrying opposiuuy

directed magnetic field in t~e interplanetary space. The

idea of HCS was introduced by Schultz (1973) to explain the

IMF sector structure and was developed later by Alfven

(1977), Svalgaard et al. (1974, 1975), Svalgaard and Wilcox

(1976b), Saito (1975) and Levy (1976). A phenoQenological

model of HCS is developed by Kaburaki and Yoshii (1979) and

Kaburaki and Imai (1983). A schematic drawing of the

heliospheric current sheet is shown in figure 1.2.

29

+JOJOJ

..c:Ul

+Jl:::OJHH::loo

.r-!HOJ

..c:~Ulo

.r-!rlOJ

..c:OJ

..c:+J

4-lotJ1l:::

.r-!~ruH'"d

o.r-!+JruEiOJ

..c:oU)

N

30

Observations of heliospheric current sheet

Potential field modeling

Schatten et al. (1969) and Altschuler and Newkirk

(1969) independently introduced the concept of a potential

field model with a spherical source surface surrounding and

concentric with -the Sun. In this model the line of sight

photospheric magnetic field is used to determine the

configuration of large Scale heliospheric magnetic field

assuming that

(1) There is no currents in the region bet\Jeen Sun and the

source surface and that

(2) at the source surface the Qagnetic field is purely

radial.

The locus of the HCS is given by connecting the points where

the radial field goes to zero at the source surface.

Pneumann et al. (1978) and Wilcox et al. (1980) utilised

this model to infer HCS structures during 1973 and 1976

respectively. Later Hoeksema (1984) refined this Qodel to

cOQpute HCS structure for the period 1976-1983 using low

resolution photospheric magnetic field observations made at

the Hilcox Solar Observatory, Stanford. Several corrections

31

are to be applied before computing heliospheric magnetic

field from line of sight photospheric magnet~c field. Some

of them are

(1) The polar field observed is to be corrected using

annual variation in polar field strength (Svalgaard et

al., 1978).

(2) Zero-effect is to be removed from the data and

(3) the source surface radii is fixed at an optiwum

distance from the Sun (_2.5 R~) where

heliospheric field correlates better with

observations of IMF near earth.

the computed

the actual

Further one needs to minimise the effects of

evolution of large scale magnetic field within a solar

rotation. Hoeksema and Scherrer (1986) extended the

computation of heliospheric magnetic field up to 1985 and

recently by Hoeksema (1989) up to 1988. Source surface

magnetic field map for the carrington rotation 1740 is ShO\ln

in figure 1.3.

~ILCDX SOLAR 08SERUATORY

SOURCE

2 0

I-c

SURFACE r'lAGNETIC FIELD 0, ± 1 ,.,'- 5, 10, 2 0 r"lICROTESL':'

<: " I

~±fi£T::W~=-

e:

o 30 60 90 1 2 0 :) 0 18(1 210 <: <1 0 " ,. (1 :;. (; (I (> .: ":, I~ I

"7 d JI" ,'./

Fig.l.3. Source surface magnetic(Hoeksema and Scherrer,

field map1986)

for the carrington rotation 1740

et al.

33

K-corona observations

The locus of the brightest regions in a synoptic

K-corona map will give the position of HCS in heliocentric

co-ordinates (Hundhausen, 1977). This method, known as

maximum brightness curve (HBC) method is utilised by Burlaga

(1981) and Bruno et al (1982, 1984) to infer HCS

geometry during 1976-1977. Bruno et al. (1984) found that

the RCS structures cornputed from K-corona observations show

generally good agreement \lith the RCS inferred fror:1

potential field modeling a photospheric magnetic field data

during 1976 except for minor differences. This result can

also be found from similar other studies (Pneurnann, 1976i

Pneumann et al., 1978i Wilcox and Rundhausen, 1983).--Korzhov (1977) has used the OSO-7 K-corona obser-

vations and synoptic K-corona observations from ground based

K-corona meters to infer RCS structures during 1971-1978

(Korzhov, 1982). Newkirk and Fisk (1985) identifies the RCS

on the centre of the bright band of coronal streamers in the

K-corona observations for the period 1965-1982. Saito and

Swinson (1986) utilised the mid-line method to infer HCS

position from K-corona observations. In their method RCS

latitude at a given longitude is given by the mid-points of

the boundary of polar coronal holes in the opposite

34

heliohemispheres. All the above investigations found a good

correlation between the predicted polarity of IMF from their

inferred HCS positions and actual observations of IMF near 1

AU. The inclination of HCS observed in K-corona is compared

with the same observed near earth's orbit by Behannon et ale

(1983).

Other methods

Yearly average solar wind maps covering nearly

+70 0 of solar latitude and entire solar longitude can be

made using IPS observations (Coles et al., 1980i Rickett and

Coles, 1983; Kojima ar:d Kakinuma, 1987). The position ot ImJ

speed solar wind belt (V<400 -1 in givekms ) these maps a

good approximation to the HC~ POS1L.LUIlS inferred using K-

corona oDservations or potential field modeling of solar

magnetic field.

Stewart (1985, 1987) developed a method to infer

RCS position using synoptic plots of solar radio noise

storms in the interval 1973-1984. The dividing line between

opposite noise-storm polarities appears to be a good

representation of the HCS up to latitudinal displacements

+50 0 from the solar equator.

In addition to these methods IMF p01arity observed

near earth can also be used to infer the properties of HCS

35

(Svalgaard and Wilcox, 1976b; Tritakis, 1984a; Villante et

al., 1979). Geomagnetic indices also provide clues to the

structure of the HCS (Simon and Legrand, 1987; Saito and

Saito, 1986; Triskova, 1988; Oksman and Kataja, 1986). The

interaction between comet tails and the HCS is also used to

infer HCS properties in the past (Neidner, 1982).

Evolution of HCS with sunspot cycle

During sunspot minimum the HCS stays very close to

the solar equator. For e.g. HCS is within +150

during

several solar rotations in 1976 (Hoeksema, 1984; Newkirk and

Fisk, 1985). But as the solar activity builds up, HCS

extends to higher solar latitudes typically >70 0. Even

though the structure of IMF sector structure observed near

earth is simple consisting of either two or four sectors per

solar rotation. HCS configuration during sunspot maximum is

complex. Multiple HCS structures are usually observed during

this period (Hoeksema, 1984; Korzhov, 1977). During solar

cycle 21 through 1978 to 1981, the latitudinal extension of

HCS was >50 0• But later after 1985, the latitudinal

extension dropped to 22.5 0 and further decreased near 1987

(Behannon et al., 1988). Thus one can find a systematic

variation of the latitudinal extension of the HCS with the

solar cycle. In addition to this, the number of warps in the

36

HCS also vary with sunspot cycle. In solar cycle 21 the

warps of HCS extended up to 50 0 during several years of

observation (Hoeksema and Scherrer, 1986). The geometry of

HCS is not observed to be symmetric about the heliographic

equator. During 1971-1977 the HCS is extended more in the

southern heliosphere (Burlaga et al., 1981; Korzhov, 1983).

Tritakis (1984a) provides evidence of the asymmetry in HCS

placement about solar equator from long term analysis" of

IMF polarity data. The complexity of IMF sector structure

observed near earth (two sector structure or four sector

structure) depend on the structure of the HCS during various

phases of the sunspot cycle (Hoeksema, 1984; Gonzalez and

Gonzalez, 1987).

Hundhausen (1977) and Thomas and Smith (1981)

assumed a tilted dipole heliomagnetic configuration to

explain the structure of the HCS. But during most of the

solar cycle period higher order solar magnetic multipoles

distort the structure of HCS from a simple sinusoidal

structure due to a dipolar solar magnetic field. Even near

sunspot minimum when the heliomagnetic field is

predominantly dipolar, the structure of the HCS in the

equatorial region of the Sun depends on a smaller

quadrupolar field (Bruno et al., 1982). During most of the

period one can approximate the large scale heliomagnetic

field with a magnetic dipole and a magnetic quadrupole

37

contributing to the structure of the HCS (Hoeksema, 1984;

Saito and Swinson, 1986; Levine, 1980).

The contribution of higher order multipoles to the

heliomagnetic field increases with sunspot activity. During

high solar active period often the quadrupolar contribution

exceed the dipolar contributions to the HCS and the octupole

contribution becomes comparable to that of dipole (Hoeksema,

1984; Balough, 1986). Near solar maximum the HCS configura­

tion is very complex and multiple HCS structures are

commonly observed during this period.

Another important feature is that the dominant

polarity of IMP above or below the HCS is controlled by the

22 year heliomagnetic cycle. During solar cycle 21 the

dominant polarity above or below the HCS reversed with the

reversal ln the polarity of solar poles during 1980

(Hoeksema, 1984; Smith et al., 1986; Saito et al., 1987;

Behannon et al., 1989).

Microstructure of the HCS

It was found from the space-craft observations of

IMP sector boundaries near 1 AU that HCS is~ non-null

magnetic field region where interesting reconnective

processes may take place (Klein and Burlaga, 1980; Behannon

and Neubauer, 1981; Behannon et al., 1981; Villante and

38

Bruno, 1982). Behannon et ale (1981) conducted a detailed

study on the fine scale characteristicS of IMF sector

boundaries observed near earth during 1974-1975. The field

configuration of the HCS is found to be associated with

directional discontinuitiesj 'thick' as well as

portions of the HCS is observed. Behannon et ale

'thin'

(1981)

observed a thickness of 3xl014 cm associated with the HCS.

Eslevich and Filippov (1988) show that the plasma and

magnetic field in the HCS does not undergo appreciable

additional change due to .interplanetary dynamics during its

transit from Sun to 1 AU. HCS structure observed at the

earths orbit is then basically determined by its projections

at the solar surface. The variation of magnetic field in the

HCS implies a rotation or tangential transition which the

magnetic vector experience as they move across the HCS

titled at angle 8 with respect to the ecliptic plane. The

rotation of magnetic vector inside the HCS is likely due to

the differential rotation of the Sun. Bruno and Bavassano

(1987) investigated on the dependence of IMF power spectrum

upon the angular displacement from HCS using Helios

observations. At the low wave numbers (K<2xlO- 6km- l ) the

magnetic field controlled by slow solar wind has spectral

characteristics different from those of fast wind. Close to

the HCS a -1spectrum as K is seen. Recently Bruno and

Bavassano (1988) studied the variation of the power of IMF

39

fluctuations \vith angular distance from heliospheric current

sheet.

Deformation of HCS due to Solar wind velocity gradients

In the highly idealised case of a totally uniform,

steady, solar wind, the shape of the HCS is independent of

the distance from the Sun. Suess and Hildner (1985) have

considered the effect of an inhomogenous solar wind with

velocity varying from point to point along the HCS near the

source surface causing a distortion in the HCS geometry at

large distances from the Sun. Suess et al. (1986)

demonstrated this effect from observed azimuthal velocity

gradient along the Hes during a solar rotation period in

1977. But the observations of the distant HCS structure by

Pioneer and Voyager missions does not show appreciable

distortions from the near earth HCS structure as predicted

by Suess and Hildner (1985) (Behannon et al., 1989; Mihalov

et al., 1989).

Relation with cosmic ray propagation and comets

Svalgaard and Wilcox (1976b) suggested that the

solar cycle variations on the cosmic ray intensity observed

near earth could be explained in terms of the change in tilt

40

of the HCS with solar cycle. Jokipii and Thomas (1981) were

able to model the effect of simple tilted dipole

configuration of HCS on the propagation of galactic cosmic

rays. Heliospheric current sheet is thus a natural defence

against intense galactic cosmic rays entering the

heliosphere. Increasing the tilt of the HCS significantly

decreases the flux of the cosmic rays at earth. During solar

cycle 21 and near the minima of solar cycle 22 the changes

in the tilt of the HCS fairly correlated with the cosmic ray

intensity variations observed at earth (vvebber and Lockwood,

1988; Hoeksema, 1984).

Neidner and Brandt (1978) suggested that the

observed disconnection events (DE) of the ionic tails of the

comets from the nucleus is associated with the interaction

of the cometary plasma with the HCS. Neidner (1982)

identified such 72 DE from observations of comet tails since

1892 and he has inferred the latitudinal extent and tilt

properties of the HCS during these events. During the recent

spacecraft encounters with comet Halley several DE were

observed in relation with the HCS (Brandt and Neidner, 1987;

Saito et al., 1987; Lundstedt and Magnusson, 1987).