La Heliosfera Sergio Dasso 1,2 1 Instituto de Astronomía y Física del Espacio (IAFE), CONICET-UBA, Argentina 2 Departamento de Física, Facultad de Ciencias Exactas y Naturales, UBA, Argentina Departamento de Física Juan José Giambiagi
Clase 3: Turbulencia en el medio interplanetario El viento solar transitorio Regiones de co-rotación Estructuras transitorias Eyecciones de masa coronal Nubes magnéticas interplanetarias
Turbulence in the SW In HD turbulence, the kinetic energy developes a cascade from large scale fluctuations to small scales through the intertial range At a proper small scale ( dissipation range ) the fluctuations start to be damped and the energy is deposited as heating of the fluid However, MHD is more complex: Turb MHD is anisotropic, B 0!!! Direction of k (respect to B 0 ) crucial to wave modes and wave-particle interaction Diff theories make diff prediction about the evolution of E b +E k (e.g., IK 65 alternative to K41) For λ<λc, w-p, new physcis can dominate fluctuations (e.g., dispersion range): e-! Polarization turns to be crucial (branches) Many other differences
Turbulencia Hay evidencia muy sólida que indica que la turbulencia está presente tanto en el medio interestelar local como en la heliosfera En el medio interestelar local, desde observaciones remotas? Armstrong et al., Nature, 1981
Turbulencia Hay evidencia muy sólida que indica que la turbulencia está presente tanto en el medio interestelar local como en la heliosfera En el medio interestelar local, desde observaciones remotas En el viento solar, desde observaciones in situ Cascada de energía magnética (Voyager at 1 AU): E~k -5/3 Transporte en escalas (rango inercial)? Armstrong et al., Nature, 1981 From a lot of authors from 70 s to today
Courtesy from Bill Matthaeus
Power spectra at intertial range: -5/3 [K41]
Turbulencia Hay evidencia muy sólida que indica que la turbulencia está presente tanto en el medio interestelar local como en la heliosfera En el medio interestelar local, desde observaciones remotas En el viento solar, desde observaciones in situ Cascada de energía magnética (Voyager at 1 AU): E~k -5/3 Transporte en escalas (rango inercial)? Armstrong et al., Nature, 1981 From a lot of authors from 70 s to today
Turbulencia Alfvénica en el viento solar Fuerte correlación entre v y b [Belcher and Davis 1971] Estructuras de Elsässer: ± z = v ± B 4πρ Acoplamiento no lineal únicamente si se activan ambos modos (ondas de Alfvén)
How Parkerian is the SW at 1AU? [From Borovsky 2010] Longitude deviated from the expected value, according with the Parker spiral
Distorted (advected) dipolar configuration The Sun rotates and there is a tilt between magnetic dipole and
Corotating Interaction Regions
Estructuras de gran escala en la heliosfera observadas con técnicas de centelleo Se puede observar la estructura de espiral: regiones de compresión y rarificación Estos datos fueron observados desde Tierra con técnicas de centelleo: Fluctuaciones de intensidad de ondas de radio que responden a n y turb en SW (Bernie Jackson, para mas información: http://cassfos02.ucsd.edu/solar/tomography/) June 23, 1994 to July 20, 1994 (Carrington Rotation 1884)
Ejective transient structures: Coronal Mass Ejections Drawing 1860 LASCO-SOHO
Ejective transient structures: Coronal Mass Ejections Drawing 1860 LASCO-SOHO (2000)
Transients are ejected from the Sun toward the SW From [Zurbuchen & Richardson, Space Science Rev, 2006] Subset of ICMEs observed as Magnetic Clouds (MCs) in the SW Observed properties: - Low Tp - Smooth and large rotation of B - Large intensity of B - Low proton plasma β p Snow thrower effect Shock waves driven by fast MCs (and a turbulent sheath of large n) Cold structures (low Tp) e - s flows along B (>100eVs): proxy of magnetic connectivity Parker spiral B Smooth and large coherent rotation of B (helical structures), B increased respect to the SW Low plasma beta (β p ) Fmag
Connectivity to the Sun Larson et al. [GRL, 1997] Estimation of the leg s length: L~1.2AU L=v elec (t f -t 0 ) t 0 :type III burst (e- beam: Lungmuir) ω e ~n 1/2 e (14MHz 10KHz) t f :in situ (e-: 20KeV) Larson et al. [GRL, 1997] Correlation between absence of e - flux and high β value When suprathermal e - fluxes are not observed... disconnection or scattering due to wave-particle interaction?
Open and closed fields in magnetic clouds Suprathermal (320 ev) electron pitch angle distributions for 8 Nov 04 open closed magnetic cloud Shodhan et al. [2000] Counterstreaming suprathermal electrons indicate field lines connected to the Sun at both ends (closed) On average, clouds are nearly half open
Local simplification: Cylindrical slide of the MC
Modeling magnetic clouds From in situ observed B possible to orient the local axis of the flux rope, to model it, and to compute the content of MHD invariants
Cross-section shape of MCs from multispacecraft observations (assuming a Grad Shafranov equilibrium) Cross-section: ~ circular (due to magnetic tension) MC reconstruction of the magnetic field from 1D data (magnetic + plasma pressure balance) ( Liu et al. 2008 )
Magnetic clouds are astrophysical objects with very low proton plasma beta ( ~ 10-1 10-2 ), so that their internal dynamics is expected to be controlled mainly by magnetic fields From Nakwacki et al., 2008
Fluxes and H from models Linear Force Free Field [Lundquist, ArkFys 1950] Taylor s state B=B z (r)z + B φ (r)φ B B = 2 0 τ 0B ( τ = cte) = B0[ J 0(2τ 0r) z + J1(2τ 0r) φ] τ(r)=dϕ/dz=b ϕ (r)/rb z (r): amount of magnetic field twist τ 0 =τ(r~0) B 0 is the magnetic field intensity at the cloud axis
MCs and ICMEs typically show expansion (larger velocity in the front and smaller in the back) From Nakwacki et al., 2008 Possible to (in situ) observe expansion along the radial direction (from the Sun) But, strong deformation of helical structure is expected if expansion is significantly different in the other directions So that, approximately isotropic expansion is expected
Different MCs observed at different solar distances D (good for studies of global expansion) S~D 0.97 Modeling evolution of MCs from assuming: (i) conservation of mass, magnetic fluxes (ii) isotropic self-similar expansion (iii) S~D Then, n p ~D -3, B~D -2 [From Kumar & Rust, JGR 1996] B~D -1.8 Large uncertainties (only a few observed events) Improvements Better determination of scale law exponents (more events): S~D 0.8±0.1 [e.g., Bothmer & Schwenn 98; Leitner et al. JGR 07; Gulisano et al., 2010] Moving boundary model [Demoulin & Dasso A&A 09]: For SW pressure decay as P sw (D)~D -np, global expansion expected as S ~ D np/4 ~ D 0.7
F y, cloud Cumulative flux F y /L L ( x). = X x in dx' B y, cloud ( x') y x X in From B=0 and invariance of B along the main axis of the cloud (valid to a general 2D-shape!): flux rope dx B y, cloud ( x) = 0
F y, cloud Cumulative flux F y /L L ( x). = X x in dx' B y, cloud ( x') y x X in From B=0 and invariance of B along the main axis of the cloud (valid to a general 2D-shape!): B y,cloud flux rope dx B y, cloud ( x) = 0 Agreement between diff. authors on start MC. However, diff. end boundaries chosen by previous authors [Lepping et al., JGR 97; Larson et al., GRL 97; Janoo et al., JGR 98] Cancelation of F y,cloud a magnetic discontinuity at [from Dasso et al., A&A 2006]
Magnetic Reconnection? Back was part of flux rope before? Back -Also found in other MCs [Dasso et al., 2007, Mostl et al., 2008, Ruffenach et al., 2012] -Reconnection is efficient in heliosphere because Hall effect increases its rate [e.g., Morales et al., JGR 2005] -Reconnection in ICMEs boundaries (from observations [Gosling et al., 2005]; from numerical simulations [Taubenschuss, 2010]) Back!! Possible explanation: The flux rope was partially pealed from its front. However, B in the back still organized, showing ~ properties of a flux rope and not of SW Erosion affects geoeffectiveness! (less amont of flux lower time range for Bs)
Fin clase 3