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THE SOLAR WIND. 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. P.K. Manoharan Radio Astronomy Centre National Centre for Radio Astrophysics

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slide1

THE SOLAR WIND

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

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, 2007

Kodai IHY School

December 10-22, 2007

slide2
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
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
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!?)
solar wind
Solar Wind
  • 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'.
slide9

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.

slide11

SUN

Geospace

solar wind2
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
heliosphere and solar wind studies exploring heliosphere in 3 d
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
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)
    • At the base of the corona, E < 0; for b = 0.3, E > 0

( b<<1 )

supersonic flow
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
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
slide22

Large spread

Flow speed (km/s)

Heliocentric distance (Rs)

after Esser et al. (1997)

slide24

bias by waves

Harmon & Coles 2005

slide26
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
Origin of slow SW(seCH)

Coronal hole origin

but

⇒ extra momentum source

in lower corona

Large NV

High To (in seCH)

Enhanced Heating

in lower corona

seCH

Strong B (in seCH)

slide30

Flux expansion ratef

Magnetic field intensity B

large scale structure of solar wind
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)
solar wind measurements
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

Interplanetary Scintillation

Radio source

L-O-S

Sun

Earth

slide36

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

slide37

Ooty IPS measurements: Density Turbulence and Speed

of the Solar Wind in the Inner heliosphere

February 25 –

March 25, 2005

CR2027

quasi stationary solar wind large scale structure and long term variations
Quasi-stationary solar windLarge-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 large-scale structure of the coronal magnetic field over the solar activity cycle.

coronal holes
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 Cycle 23 – Solar Wind Density Distribution

Solar Wind Density Turbulence (Ooty)

radial evolution of cirs
Radial Evolution of CIRs

75 solar radii

100 solar radii

expansion

150 solar radii

radial evolution of cmes
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
slide47

June 25, 1992

West Limb CME on June 25, 1992* X3.9 Flare, X-ray LDE

Manoharan et al. ApJ., 2000

Type-IV

cme in the interplanetary medium
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
CME Propagation Speed (from Sun to Earth)

Height – Time plot

Radial Evolution of Speed

K.E. lost/dissipated within <100Rsun

~1032 erg

VCME ~ R-0.08 at R < 100 Rsun

VCME ~ R-0.72 at R > 100 Rsun

slide53

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
  • 2-step effective acceleration
  • Transition around 70 – 80 R
  • at R < 70 R: -0.3 < β < +0.06
  • at R > 70 R: -0.76 < α < 0.58
    • slope > 0 : acceleration
    • slope < 0 : deceleration
  • index ‘β’ shows no significant dependence
  • on the initial speed of the CME
  • index ‘α’ shows dependence on the initial
  • speed

Speed Profiles: VCME(R)

deceleration

constant speed

acceleration

Manoharan 2006

slide54

CME on December 13, 2006

g-index

Shock

CME

g-index

Speed

g-index

Shock

CME

Speed

Speed

slide55

|B| (nT)

Bz (nT)

V (km/s)

N

T (K)

Pressure

slide58

Cosmic ray precursors of the CME arrival at Earth

Observation the network of neutron monitors.

Yellow circles : excess, Red circles : deficit

slide59

CRs from FD region travel to the upstream Earth with the speed of light overtaking the shock ahead.

Munakata et al., JGR, 105, 2000

slide60

RL

Sun

We deduce(t) from the observed (t) & B(t)

(- (t) points toward the flux rope center)

Munakata et al., ASR, 2005

slide61

Geometry of magnetic flux rope in Halloween CME

from Cosmic Ray data

from ACE IMF data

Kuwabara et al., JGR, 31, L19803, 2004

slide62

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
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.
slide66

Ooty Radio Telescope (ORT)

Various Astronomical Studies

  • 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)

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
Four-station system for IPS

102km

126km

98km

109km

131km

slide68

Cross correlation

0

Lag time

Multi-station IPS observations

  • 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
    • Baseline should be parallel to the projected solar wind flow direction.
slide70

Point Source, Θ ~ 15 mas

Strong

scintillation

Weak scintillation

Scintillation index – Heliocentric Distance Plots

solar wind density turbulence also spectrum
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

solar wind speed
Solar wind Speed
  • Solar wind speed and Density turbulence spectrum, ΦNe(q)
    • By suitably transforming and calibrating the intensity scintillation time series
slide83

α=3.0

Φ ~ q-α

Solar wind speed

Power-law index

α=3.9

Effects of power-law index, solar wind speed,

And source size

Compact source size

slide86

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:

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%)

acceleration zone

‘zero’ acceleration line

initial speed arrival time at 1 au
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

slide88

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)

slide89

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).