p.k. manoharan radio astronomy centre national centre for radio astrophysics tata institute of...

P.K. ManoharanRadio Astronomy CentreNational Centre for Radio AstrophysicsTata Institute of Fundamental ResearchOoty 643001, India

mano@ncra.tifr.res.in

P.K. ManoharanRadio Astronomy CentreNational Centre for Radio AstrophysicsTata Institute of Fundamental ResearchOoty 643001, India

mano@ncra.tifr.res.in

THE SOLAR WINDTHE SOLAR WIND

Kodai IHY SchoolDecember 10-22, 2007Kodai IHY SchoolDecember 10-22, 2007

• 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

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'.

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.

SUNSUN

GeospaceGeospace

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)

– At the base of the corona, E < 0; for b = 0.3, E > 0

( b<<1 )

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

Flo

w s

peed

(km

/s)

Heliocentric distance (Rs) after Esser et al. (1997)

Large spread

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)

High To (in seCH)

Strong 　 B (in seCH)

⇒ extra momentum source in lower corona

Enhanced Heatingin lower corona

Coronal hole origin

Large NV but

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

Sun

Earth

Radio source L-O-S

Computer Assisted Tomography analysis

can remove the line-of-sight integration imposed on the solar wind parameters also provides high spatial resolution

Solar rotation and radial outward flow of the solar wind provide the 3-d structure of the solar wind at different view angles

Ooty IPS measurements: Density Turbulence and Speed of the Solar Wind in the Inner heliosphere

February 25 –March 25, 2005

CR2027

1999

20001991

Solar Cycle Dependence

Quasi-stationary solar windLarge-scale structure and long-term variations

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.

Constant level of electron density fluctuations (Ne), observed using the Ooty Radio Telescope, during minimum and maximumof 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

150 solar radii

75 solar radii

100 solar radii

expansion

Solar Cycle 23 – Solar wind Speed Distribution

IPS Imaging of interplanetary disturbances (CIRs and CMEs)

Sun

Earth

CME

Radio Source

Shock

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

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

Manoharan et al. ApJ., 2000

June 25, 1992

Type-IV

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 Images50 - 250 Rsun

Microsoft　PowerPoint　

Presentation

CME Propagation Speed (from Sun to Earth)

Height – Time plot Radial Evolution of Speed

VCME ~ R-0.08 at R < 100 Rsun

VCME ~ R-0.72 at R > 100 Rsun

K.E. lost/dissipated within 100Rsun

~1032 erg

A fast CME EventJanuary 20, 2005

LASCOGopalswamy et al. 2005

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)

acceleration

constant speed

deceleration

Manoharan 2006

Speed

Speed Speed

g-index

g-index

g-index

Shock

Shock CME

CME

CME on December 13, 2006

V (km/s)

|B| (nT)

Bz (nT)

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

cut-off (inertial) scale = VA/P

= N–1/2

VA Alfven speedP Proton cyclotron frequencyN Plasma density

Density turbulence spectrum associatedwith propagating CME

• 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.

Summary

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)

High-sensitivity IPS measurement using Ooty Radio Telescope provide

– Speed of the solar wind– Density turbulence spectrum Operated by

Radio Astronomy CentreNational Centre for Radio AstrophysicsTata Institute of Fundamental Research(NCRA-TIFR)Ooty, India

Giant Meter wavelength Radio Telescope (near Pune)Multi-frequency synthesis imaging system27-km baseline 30 antennas of each 45 m diameter

Various Astronomical Studies

Four-station system for IPS

102k

m

126k

m

131km

98km

109k

m

Lag time0

Cross correlation

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.

2/1

2

2

I

(t) ΔI

source theofintensity mean

nsfluctuatiointensity of rmsm

I(t)I(t)ΔI

Scintillation Index (m)

Point Source, Θ ~ 15 mas

Scintillation index – Heliocentric Distance Plots

Weak scintillation

Strongscintillation

Multi-frequency IPS

source

Earth

2e

2 (R)dzΔN m

R ΔN 0.4 4.42e

Radial dependence of density turbulence

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. theambient wind identifies the presenceof the CME along the line-of-sight direction to the radio source

t)tr,ΔI(r )t,ΔI(rtr, ρ oooo

t)dtf i2πexp( t)ρ(0,2π

1PI

(f)dfP

I

1m I2

2

IPS – Power Spectrum

Solar wind Speed

Solar wind speed and Density turbulence spectrum, ΦNe(q)

• By suitably transforming and calibrating the intensity scintillation time series

αβ22

2y

z

p

2e qRθ)z,V(q,

2k

zq4sindq

(z)V

dzλ)r (2π P(f)

IPS temporal power spectrum

q-α

Power-law index Solar wind speed

Compact source size

Φ ~ q-α α=3.0

α=3.9

Effects of power-law index, solar wind speed,And source size

Solar Wind Density Turbulence and Speed (3 days)

CME Initial Speed vs Acceleration Slope at R > 70 R

α = 0.2-6.4×10-4V+1.1×10-7V2

‘zero’ acceleration line

acceleration zone

deceleration zone

V = 380 km/s

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%)

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

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)

“Interplanetary Scintillations” (IPS)

intensity fluctuations caused by the solar wind density turbulence

This time series transformation provides the temporal power spectrum

Density Turbulence Spectrum

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