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P.K. Manoharan Radio Astronomy Centre National Centre for Radio Astrophysics Tata Institute of Fundamental Research Ooty 643001, India

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Presentation on theme: "P.K. Manoharan Radio Astronomy Centre National Centre for Radio Astrophysics Tata Institute of Fundamental Research Ooty 643001, India"— Presentation transcript:

1 P.K. Manoharan Radio Astronomy Centre National Centre for Radio Astrophysics Tata Institute of Fundamental Research Ooty 643001, India P.K. Manoharan Radio Astronomy Centre National Centre for Radio Astrophysics Tata Institute of Fundamental Research Ooty 643001, India THE SOLAR WIND Kodai IHY School December 10-22, 2007 Kodai IHY School December 10-22, 2007

2 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.

3 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)

4 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 – 5x10 5 K –Huge convection cells lead to jet-like phenomena Corona –extends from chromosphere to several R  –extremely hot, 3x10 6 K (causes high state of ionization) –energy transport by magnetic fields (heating!?)

5 X-ray Corona

6 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'.

7 LASCO Observation – Comets and Coronal Mass Ejections

8 Solar Wind

9 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.


11 SUN Geospace



14 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

15 Hourly average of solar wind speed. density and thermal speed measured at 1 AU

16 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


18 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 ( b<<1 )

19 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

20 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

21 Suzuki, ApJ 2006

22 Flow speed (km/s) Heliocentric distance (Rs) after Esser et al. (1997) Large spread

23 Axford et al.

24 bias by waves Harmon & Coles 2005

25 High-Speed Solar Wind - Coronal Hole Region

26 When a polar coronal hole shrinks to small size at the solar maximum, it becomes the source of slow wind.

27 Origin of slow SW (seCH) High To (in seCH) Strong B (in seCH) ⇒ extra momentum source in lower corona Enhanced Heating in lower corona Coronal hole origin Large NV but seCH

28 after Kojima et al., 1999


30 Flux expansion rate f Magnetic field intensity B

31 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)

32 High- and Low-Speed Solar Wind

33 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


35 Interplanetary Scintillation Sun Earth Radio source L-O-S

36 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

37 Ooty IPS measurements: Density Turbulence and Speed of the Solar Wind in the Inner heliosphere February 25 – March 25, 2005 CR2027


39 1999 20001991 Solar Cycle Dependence

40 Quasi-stationary solar wind Large-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 (  N e ), observed using the Ooty Radio Telescope, during minimum and maximum of solar activity cycle.

41 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

42 Solar Cycle 23 – Solar Wind Density Distribution Solar Wind Density Turbulence (Ooty)

43 Radial Evolution of CIRs 150 solar radii 75 solar radii 100 solar radii expansion

44 Solar Cycle 23 – Solar wind Speed Distribution

45 IPS Imaging of interplanetary disturbances (CIRs and CMEs) Sun Earth CME Radio Source Shock

46 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

47 West Limb CME on June 25, 1992 * X3.9 Flare, X-ray LDE Manoharan et al. ApJ., 2000 June 25, 1992 Type-IV

48 Some example of November 2003 CMES

49 Fast CME on April 2, 2001: Ooty Images

50 CME in the interplanetary medium LASCO Images <30 Rsun Waves Radio Spectrum Ooty Scintillation Images 50 - 250 Rsun

51 CME Propagation Speed (from Sun to Earth) Height – Time plotRadial Evolution of Speed V CME ~ R -0.08 at R < 100 Rsun V CME ~ R -0.72 at R > 100 Rsun K.E. lost/dissipated within  100Rsun ~10 32 erg

52 A fast CME Event January 20, 2005 LASCO Gopalswamy et al. 2005

53 V CME (R) of 30 CMEs IPS & LASCO provide sky-plane speeds Include constant speed, accelerating and decelerating events V CME (R) can be represented by power-law forms: V CME (R) ~ R -β R < 50 R  V CME (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: V CME (R) acceleration constant speed deceleration Manoharan 2006

54 Speed g-index Shock CME CME on December 13, 2006

55 V (km/s) |B| (nT) Bz (nT) N T (K) Pressure


57 Neutron Monitor Station Count Rates

58 Cosmic ray precursors of the CME arrival at Earth Observation the network of neutron monitors. Yellow circles : excess, Red circles : deficit

59 CRs from FD region travel to the upstream Earth with the speed of light overtaking the shock ahead. Munakata et al., JGR, 105, 2000

60 RLRL Sun We deduce   (t) from the observed  (t) & B (t) (-   (t) points toward the flux rope center) Munakata et al., ASR, 2005

61 Geometry of magnetic flux rope in Halloween CME from Cosmic Ray data from ACE IMF data Kuwabara et al., JGR, 31, L19803, 2004

62 Spectra associated with ambient low- and high-speed solar wind flows Solar wind Density turbulence spectrum cut-off (inertial) scale = V A /  P = N –1/2 V A Alfven speed  P Proton cyclotron frequency N Plasma density Density turbulence spectrum associated with propagating CME

63 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


65 Thank you

66 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, ± 60 o ) High-sensitivity IPS measurement using Ooty Radio Telescope provide –Speed of the solar wind –Density turbulence spectrum Operated by Radio Astronomy Centre National Centre for Radio Astrophysics Tata Institute of Fundamental Research (NCRA-TIFR) Ooty, India Giant Meter wavelength Radio Telescope (near Pune) Multi-frequency synthesis imaging system 27-km baseline 30 antennas of each 45 m diameter Various Astronomical Studies

67 Four-station system for IPS 102km 126km 131km 98km 109km

68 Lag time 0 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.

69 Scintillation Index (m)

70 Point Source, Θ ~ 15 mas Scintillation index – Heliocentric Distance Plots Weak scintillation Strong scintillation

71 Multi-frequency IPS



74 Radial dependence of density turbulence

75 Solar wind Density Turbulence (also spectrum) Density Turbulence * Scintillation index, m, is a measure of level of turbulence * Normalized Scintillation index, g = 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


77 IPS – Power Spectrum


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

80 IPS temporal power spectrum


82 q-αq-α

83 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


85 Solar Wind Density Turbulence and Speed (3 days)

86 CME Initial Speed vs Acceleration Slope at R > 70 R  α = 0.2-6.4×10 -4 V+1.1×10 -7 V 2 ‘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 [~ (V CME – V AMB ) 2 ] shows good linear correlation (~97%)

87 Initial Speed – Arrival Time at 1 AU T CME = 109 - 0.5 × 10 -1 V CME + 1.1 × 10 -5 V 2 CME hours V CME = 400 km/s, T CME = 90 hours (considerable assistance by CME expansion) V CME = 2000 km/s, T CME dominated by interaction Includes energy provided by CME Expansion + SW interaction

88 is wavelength of observation; r e is classical electron radius. F diff (q) = Fresnel diffraction filter (attenuates low-frequency part of the spectrum) F Source (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

89 Axial Ratio of Irregularity When the density irregularities are field aligned and approximated with an ellipsoidal symmetry, the spatial spectrum of density fluctuations, Φ Ne (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. q i cut-off scale i.e., inner-scale size).

90 Thank You

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