Post on 25-Mar-2018
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).