THE SOLAR WIND
P.K. Manoharan
Radio Astronomy Centre
National Centre for Radio Astrophysics
Tata Institute of Fundamental Research
Ooty 643001, India
mano@ncra.tifr.res.in
Kodai IHY School
December 10-22,
10-22, 2007
December
• J. L. Kohl and S. R. Cranmer (eds.), Coronal Holes and
Solar Wind Acceleration}, Kluwer Academic Publishers,
1999.
• E. Marsch, Living Review in Solar Physics, vol. 3, 2006.
• M. K. Bird and P. Edenhofer,
Physics of the Inner Heliospher - I,
eds. R. Schwenn and E. Marsch,
Springer--Verlag, Berlin, 1990.
Outline
• Introduction – Solar Atmosphere
• Solar Wind
– formation
– acceleration
• Interplanetary Magnetic field
– magnetic storms
• Solar wind measuring techniques
– direct (in situ) measurements
– remote-sensing techniques
• Interplanetary Scintillation
– Speed and density turbulence
• Quasi-stationary (Steady-state) solar wind
• Transients in the solar wind (CIRs and CMEs)
Solar Atmosphere
• Photosphere
–
–
–
–
thin layer of low-density gas
allows visible photons to escape into space
currents of rising from beneath cause formation granulation
magnetic fields threading outward
• magnetic structures (sunspots, active regions, etc.)
• Chromosphere
– 3000 – 5000 km thick, above photosphere
– 5000 – 5x105 K
– Huge convection cells lead to jet-like phenomena
• Corona
– extends from chromosphere to several R
– extremely hot, 3x106 K (causes high state of ionization)
– energy transport by magnetic fields (heating!?)
X-ray Corona
• The concept of continuous flow of solar wind was developed
in 1950's
• Biermann (1951, 1957) observed comet tails as they passed
close to the Sun, and explained the formation of the tail and
its deflection by a continuous flux of protons from the Sun.
• Parker (1964) postulated the continuous expansion of the
solar corona, i.e., the outward streaming coronal gas the
'solar wind'.
Solar Wind
LASCO Observation – Comets and Coronal Mass Ejections
Solar Wind
Interplanetary Magnetic Field
Radial outflow and solar rotation – frozen-in magnetic field is dragged,
Interplanetary Magnetic Field (IMF). Coronal magnetic field and IMF
properties are intimately related.
SUN
Geospace
Solar Wind
• Supersonic outflow of plasma from the Sun's corona to IP medium
• Composed of approximately equal numbers of ions and electrons
• Ion component consists predominantly of protons (95%), with a small
amount of doubly ionized helium and trace amounts of heavier ions
• Embedded in the out flowing solar wind plasma is a weak magnetic
field known as the interplanetary magnetic field
• Solar wind varies in density, velocity, temperature, and magnetic field
properties with
–
–
–
–
–
solar cycle
heliographic latitude
heliocentric distance, and
rotational period
Also varies in response to shocks, waves, and turbulence that perturb the
interplanetary flow.
• Average values of solar wind parameters near the Earth (1 AU)
– Velocity 468 km/s
– Density = 8.7 protons/cc
– magnetic field strength = 6.7 nT
Hourly average of solar wind speed. density and thermal speed measured
at 1 AU
Heliosphere and solar wind studies
Exploring Heliosphere in 3-D
Determination of overall morphology of the Heliosphere
•
•
•
•
•
Acceleration of solar wind
Generation of high speed streams with correct V, N, and T
Coronal propagation of solar energetic particles
CME trajectory
Large-scale variation of solar wind and magnetic field and the
behavior of their turbulence levels
Formation of the Solar Wind
• For a steady state of the spherically symmetric flow
of solar wind,
– Equation of motion
– Equation of continuity
– Energy equation
– Temperature variation with distance (Parker 1964)
( b<<1 )
– At the base of the corona, E < 0; for b = 0.3, E > 0
Supersonic Flow
• at the base of the corona,
–
–
–
–
E is negative
system is stable
gravitation potential decreases as 1/r
thermal energy is governed by T(r), which a weak
function of distance, r
– for b ~ 0.3, E > 0 at R ~ 10 Rsun
– solar wind flows with supersonic speed
– gravity aids the nozzle flow (like a rocket jet)
• to explain the solar wind speed near the Sun and
in the entire heliosphere
Thermal and Wave driven
• Solar wind driven by thermal conduction
– not adequate to explain high-speeds at 1 AU
– some other non-thermal processes must play a role
– additional energy
• work done on the plasma or by heating, or both
– spectral broadening suggest substantial increase in
turbulence at the low corona (Alfven waves)
• model should address heating (ion and electron) and
damping/dissipation of waves
– at what height energy is added to accelerate solar wind
Suzuki, ApJ 2006
Large spread
Heliocentric distance (Rs)
after Esser et al. (1997)
Axford et al.
bias by waves
Harmon & Coles 2005
High-Speed Solar Wind - Coronal Hole Region
When a polar coronal hole shrinks to small size at the
solar maximum, it becomes the source of slow wind.
Origin of slow SW(seCH)
Coronal hole origin
but
Large NV
⇒ extra momentum source
in lower corona
High To (in seCH)
seCH
Enhanced Heating
in lower corona
Strong B (in seCH)
after Kojima et al., 1999
after Kojima et al., 1999
Flux expansion rate ｆ
Magnetic field intensity B
Large-scale structure of Solar Wind
• Steady-state solar wind (origin & acceleration)
– Low-speed solar wind
– High-speed solar wind (associated with coronal holes
• Disturbed solar wind (due to solar transients
generated by interactions, flares, and coronal
mass ejections)
High- and Low-Speed Solar Wind
Solar Wind Measurements
Solar wind measuring techniques
• Near the orbit of the Earth (~1 AU), the solar wind
properties are from in situ measurements
– Helios satellite measure up to ~0.3 AU
– Ulysses first spacecraft probed the polar region
• Scattering techniques provide the three-dimensional
view of the heliosphere
– various distances
– all latitudes
– long-term variations and large-scale structure of the solar
wind
Interplanetary Scintillation
Radio source
L-O-S
Sun
Earth
Solar rotation and radial outward flow of the solar wind provide
the 3-d structure of the solar wind at different view angles
Computer Assisted Tomography analysis
can remove the line-of-sight integration imposed on the
solar wind parameters also provides high spatial resolution
Ooty IPS measurements: Density Turbulence and Speed
of the Solar Wind in the Inner heliosphere
CR2027
February 25 –
March 25, 2005
Solar Cycle Dependence
1999
1991
2000
Quasi-stationary solar wind
Large-scale structure and long-term variations
Constant level of electron density fluctuations (Ne), observed
using the Ooty Radio Telescope, during minimum and maximum
of solar activity cycle.
Latitudinal variations of solar wind speed, observed using
the Ooty Radio Telescope, reveal the changes in the largescale structure of the coronal magnetic field over the solar
activity cycle.
Coronal Holes
• Significantly lower density and temperature than
the typical background corona
• Areas of the Sun that are magnetically open to
interplanetary space
– Configuration is divergent
• Observed in X-ray, EUV and radio wavelengths
that originate in the corona
• Grouped into 3 categories: polar, non-polar
(isolated) and transient coronal holes
• Sources of high-speed solar wind streams
– Give rise to recurrent geomagnetic storms
– Important in heliospheric and space weather studies
Solar Cycle 23 – Solar Wind Density Distribution
Solar Wind Density Turbulence (Ooty)
Radial Evolution of CIRs
75 solar radii
100 solar radii
expansion
150 solar radii
Solar Cycle 23 – Solar wind Speed Distribution
IPS Imaging of interplanetary disturbances (CIRs and CMEs)
Shock
Radio Source
Sun
CME
Earth
Radial Evolution of CMEs
– LASCO and IPS measurements between Sun
and 1 AU
– Halo and Partial Halo CMEs
– ICME at 1 AU (Wind and ACE data)
– Initial Speeds in the range 250 – 2600 km/s
June 25, 1992
West Limb CME on June 25, 1992
* X3.9 Flare, X-ray LDE
Type-IV
Manoharan et al. ApJ., 2000
Some example of November 2003 CMES
Fast CME on April 2, 2001: Ooty Images
CME in the interplanetary medium
LASCO Images
<30 Rsun
Waves Radio Spectrum
Ooty Scintillation Images
50 - 250 Rsun
CME Propagation Speed (from Sun to Earth)
Height – Time plot
Radial Evolution of Speed
K.E. lost/dissipated within <100Rsun
VCME ~ R-0.08 at R < 100 Rsun
~1032 erg
VCME ~ R-0.72 at R > 100 Rsun
Gopalswamy et al. 2005
LASCO
A fast CME Event
January 20, 2005
Speed Profiles: VCME(R)
VCME(R) of 30 CMEs
• IPS & LASCO provide sky-plane speeds
• Include constant speed, accelerating and
decelerating events
• VCME(R) can be represented by power-law
forms:
VCME(R) ~ R-β R < 50 R
VCME(R) ~ R-α R ~ 100 - 200 R
deceleration
• 2-step effective acceleration
• Transition around 70 – 80 R
• at R < 70 R: -0.3 < β < +0.06
constant speed
• at R > 70 R: -0.76 < α < 0.58
• slope > 0 : acceleration
• slope < 0 : deceleration
acceleration
• index ‘β’ shows no significant dependence
on the initial speed of the CME
• index ‘α’ shows dependence on the initial
speed
Manoharan 2006
CME on December 13, 2006
g-index
Shock
CME
g-index
Speed
Shock
g-index
CME
Speed
Speed
|B| (nT)
Bz (nT)
V (km/s)
N
T (K)
Pressure
Neutron Monitor Station Count Rates
Cosmic ray precursors of the CME arrival at Earth
Observation the network of neutron monitors.
Yellow circles : excess, Red circles : deficit
CRs from FD region travel to the upstream Earth with the speed of
light overtaking the shock ahead.
Munakata et al., JGR, 105, 2000
RL
Sun
We deduce (t) from the observed (t) & B(t)
(- (t) points toward the flux rope center)
Munakata et al., ASR, 2005
Geometry of magnetic flux rope in Halloween CME
from Cosmic Ray data
from ACE IMF data
Kuwabara et al., JGR, 31, L19803, 2004
Spectra associated with ambient
low- and high-speed solar wind flows
Solar wind
Density turbulence spectrum
Density turbulence spectrum associated
with propagating CME
cut-off (inertial) scale = VA/P
= N–1/2
VA Alfven speed
P Proton cyclotron frequency
N Plasma density
Summary
• CME Speed profile, V(R), shows dependence on initial speed
• CME goes through continuous changes, which depend on its
interaction with the surrounding solar wind
• Arrival time and Speed of the CME at 1 AU predicted by the
speed profile are in good agreement with measured values
• Mean travel time curve for different initial speeds suggests that
up to a distance of ~80 Rsun, the internal energy of the CME (or
its expansion) dominates and however, at larger distances, the
CME's interaction with the solar wind appears to control the
propagation
• Most of the CMEs tend to attain the speed of the ambient flow
at 1 AU or further out.
• These results are useful to quantify the ‘drag force’ imposed on
the CME by the interaction with the surrounding solar wind
and it is essential in modeling the CME propagation.
Thank you
Ooty Radio Telescope (ORT)
•
Latitude: 11°23’ North Longitude:
76°40’ East
•
Equatorially mounted, off-axis
parabolic cylinder
•
530m (N-S) x 30m (E-W)
•
Reflecting surface made of 1100
stainless steel wires
•
Feed – 1056 λ/2 dipoles
•
E-W Tracking and N-S Steering of
ORT (~9.5 hours, ± 60o)
Various Astronomical Studies
High-sensitivity IPS measurement
using Ooty Radio Telescope provide
– Speed of the solar wind
– Density turbulence spectrum
Giant Meter wavelength Radio Telescope (near Pune)
Multi-frequency synthesis imaging system
27-km baseline
30 antennas of each 45 m diameter
Operated by
Radio Astronomy Centre
National Centre for Radio Astrophysics
Tata Institute of Fundamental Research
(NCRA-TIFR)
Ooty, India
Four-station system for IPS
Multi-station IPS observations
Cross correlation
Speed of the solar wind can be computed from the
cross-correlation delay. But, it is restricted to :
• Baseline length has to be a few times longer than
the Fresnel radius, and
0
Lagto
time
• Baseline should be parallel
the projected solar
wind flow direction.
Scintillation Index (m)
ΔI  I(t)  I(t)
 rms of intensity fluctuatio ns 

m  
 mean intensity of the source 
 ΔI (t)
 
2
 I
2
1/ 2



Point Source, Θ ~ 15 mas
Strong
scintillation
Weak scintillation
Scintillation index – Heliocentric Distance Plots
Multi-frequency IPS
Radial dependence of density turbulence
m 
2
source
Earth
ΔN (R)dz
ΔN  R
2
e
2
e
4.4  0.4
Solar wind Density Turbulence
(also spectrum)
Density Turbulence
* Scintillation index, m, is a measure of level of turbulence
* Normalized Scintillation index, g = m(R) / <m(R)>
* Quasi-stationary and transient/disturbed solar wind
• g > 1  enhancement in Ne
• g  1  ambient level of Ne
• g < 1  rarefaction in Ne
Scintillation enhancement w.r.t. the
ambient wind identifies the presence
of the CME along the line-of-sight
direction to the radio source
IPS – Power Spectrum
ρ r, t   ΔI(ro , t o ) ΔI(ro  r, t o  t)
1 
PI 
ρ(0, t) exp( i2π f t)dt

2π 
1  
m  2  PI (f)df 

I  
2
Solar wind Speed
Solar wind speed and Density turbulence spectrum, ΦNe(q)
• By suitably transforming and calibrating the intensity
scintillation time series
IPS temporal power spectrum
P(f)  (2π r e λ)
2

z

dz
Vp (z)
2

q
z
2
2
β  α


 dq y 4sin  2k  V(q, z, θ) R q

 q 2z 
Φφ (q)
ΦΔI  4 sin 
2k


2
q-α
Φ ~ q-α
α=3.0
Solar wind speed
Power-law index
α=3.9
Effects of power-law index, solar wind speed,
And source size
Compact source size
Solar Wind Density Turbulence and Speed (3 days)
CME Initial Speed vs Acceleration Slope at R > 70 R
deceleration zone
V = 380 km/s
α = 0.2-6.4×10-4V+1.1×10-7V2
Aerodynamic drag force:
acceleration zone
‘zero’ acceleration line
Interaction between the CME cloud
and the ambient solar wind plays an
important role in the propagation of
CMEs
K.E. utilized/gained times α against the
“drag force” imposed by the ambient
solar wind [~ (VCME – VAMB)2] shows
good linear correlation (~97%)
Initial Speed – Arrival Time at 1 AU
TCME = 109 - 0.5 × 10-1 VCME + 1.1 × 10-5 V2CME hours
VCME = 400 km/s, TCME = 90 hours (considerable assistance by CME expansion)
VCME = 2000 km/s, TCME dominated by interaction
Includes energy provided by CME Expansion + SW interaction
Density Turbulence Spectrum
“Interplanetary Scintillations” (IPS)
intensity fluctuations caused by the solar wind density
turbulence
This time series transformation provides the temporal
power spectrum
 is wavelength of observation; re is classical electron radius.
Fdiff(q) = Fresnel diffraction filter (attenuates low-frequency part of the spectrum)
FSource(q) = Brightness distribution of the source (attenuates high frequency part)
Axial Ratio of Irregularity
When the density irregularities are field aligned and
approximated with an ellipsoidal symmetry, the
spatial spectrum of density fluctuations, ΦNe(q), for a
radio source with the finite size, θ, will be
AR is the ratio of major to minor axes (axial ratio),
which is the measure of degree of anisotropy of
irregularities (α power-law index. qi cut-off scale i.e.,
inner-scale size).
Thank You