Solar Wind Origin & Heating 1 Steven R. Cranmer Harvard-Smithsonian Center for Astrophysics

Everything is online at: http://www.cfa.harvard.edu/~scranmer/NSO/ Logistics Part 1: Observations Part 2: Theory 1. Background & history 2. In situ solar wind 3. Radio scintillations 4. Coronal remote-sensing (empirical connections between corona & solar wind) 5. Chromosphere & photosphere 6. Future instrumentation 1. Photosphere: tip of the iceberg of the convection zone 2. Chromosphere: waves start to propagate and bump into the magnetic field 3. Corona: magnetic field is king; heating still a “problem” 4. “Other” wind acceleration ideas; evolution of waves & turbulence 5. Future directions for theory...? (discussions afterward?) Everything is online at: http://www.cfa.harvard.edu/~scranmer/NSO/

Motivations Solar corona & wind: “Space weather” can affect satellites, power grids, and astronaut safety. The Sun’s mass-loss history may have impacted planetary formation and atmospheric erosion. The Sun is a “laboratory without walls” for many basic processes in physics, at regimes (T, P) inaccessible on Earth! plasma physics nuclear physics non-equilibrium thermodynamics electromagnetic theory

The extended solar atmosphere . . . Heating is everywhere . . . . . . and everything is in motion

The Sun’s outer atmosphere The solar photosphere radiates like a blackbody; its spectrum gives T, and dark “Fraunhofer lines” reveal its chemical composition. Total eclipses let us see the vibrant outer solar corona: but what is it? 1870s: spectrographs pointed at corona: Is there a new element (“coronium?”) 1930s: Lines identified as highly ionized ions: Ca+12 , Fe+9 to Fe+13 it’s hot! Fraunhofer lines (not moon-related!) unknown bright lines

The solar wind: hints 1850–1950: Evidence slowly builds for outflowing magnetized plasma in the solar system: solar flares  aurora, telegraph snafus, geomagnetic “storms” comet ion tails point anti-sunward (no matter comet’s motion) “Celestial battery:” Transcript between Portland and Boston telegraph stations (1859): Portland: “Please cut off your battery; let us see if we can work with the auroral current alone.” Boston: “I have already done so! How do you receive my writing?” Portland: “Very well indeed - much better than with batteries.”

The solar wind: prediction 1958: Gene Parker proposed that the hot corona provides enough gas pressure to counteract gravity and accelerate a “solar wind.” Momentum conservation: To sustain a wind, /t = 0, and RHS must be naturally “tuned:”

In situ solar wind: properties Mariner 2 (1962): first direct confirmation of continuous fast & slow solar wind. Uncertainties about which type is “ambient” persisted because measurements were limited to the ecliptic plane . . . Ulysses left the ecliptic; provided 3D view of the wind’s source regions. speed (km/s) Tp (105 K) Te (105 K) Tion / Tp O7+/O6+, Mg/O 600–800 2.4 1.0 > mion/mp low 300–500 0.4 1.3 < mion/mp high fast slow By ~1990, it was clear the fast wind needs something besides gas pressure to accelerate so fast!

Ulysses’ view over the poles

Exploring the solar wind (1970s to present) Space probes have pushed out the boundaries of the “known” solar wind . . . Helios 1 & 2: “inner” solar wind (Earth to Mercury) Ulysses: “outer” solar wind (Earth to Jupiter, also flew over N/S poles) Voyager 1 & 2: far out past Pluto: recently passed the boundary between the solar wind and the interstellar medium CLUSTER: multiple spacecraft probe time and space variations simultaneously

Particles are not in “thermal equilibrium” …especially in the high-speed wind. mag. field WIND at 1 AU (Steinberg et al. 1996) Helios at 0.3 AU (e.g., Marsch et al. 1982) WIND at 1 AU (Collier et al. 1996)

Fluctuations & turbulence Fourier transform of B(t), v(t), etc., into frequency: How much of the “power” is due to spacecraft flying through flux tubes rooted on the Sun? f -1 “energy containing range” f -5/3 “inertial range” The inertial range is a “pipeline” for transporting magnetic energy from the large scales to the small scales, where dissipation can occur. Magnetic Power f -3 “dissipation range” few hours 0.5 Hz

Radio scintillations v << VA v ≈ vwind ? The corona scatters radio waves from either natural or “unnatural” sources. Refractive index of the corona is perturbed by fluctuations in density and magnetic field. Speed of irregularities can be measured: v << VA v ≈ vwind ? (Spangler et al. 2002; Harmon & Coles 2005)

Overview of coronal observations Plasma at 106 K emits most of its spectrum in the UV and X-ray . . . Coronal hole (open) “Quiet” regions Active regions

Solar wind: connectivity to the corona High-speed wind: strong connections to the largest coronal holes hole/streamer boundary (streamer “edge”) streamer plasma sheet (“cusp/stalk”) small coronal holes active regions Low-speed wind: still no agreement on the full range of coronal sources: Wang et al. (2000)

Coronal magnetic fields Coronal B is notoriously difficult to measure . . . Potential field source surface (PFSS) models have been successful in reproducing observed structures and mapping between Sun & in situ. Wang & Sheeley (1990) found that flux tubes that expand more (from Sun to SS) have lower wind speeds at 1 AU. Wind Speed Expansion Factor (e.g., Arge & Pizzo 2000)

The SOHO mission SOHO (the Solar and Heliospheric Observatory) was launched in Dec. 1995 with the goal of solving long-standing mysteries about the Sun. 12 instruments on SOHO probe: solar interior (via “seismology”) solar atmosphere (images, movies, spectra) solar wind (collect particles, measure fields) interstellar gas (some light bounces back) L1 orbit provides 24-hour viewing!

High-resolution UV images of the low corona

On-disk vs. off-limb observations On-disk measurements help reveal basal coronal heating & lower boundary conditions for solar wind. Off-limb measurements (in the solar wind “acceleration region” ) allow dynamic non-equilibrium plasma states to be followed as the asymptotic conditions at 1 AU are gradually established. Occultation is required because extended corona is 5 to 10 orders of magnitude less bright than the disk! Spectroscopy provides detailed plasma diagnostics that imaging alone cannot.

The UVCS instrument on SOHO 1979–1995: Rocket flights and Shuttle-deployed Spartan 201 laid groundwork. 1996–present: The Ultraviolet Coronagraph Spectrometer (UVCS) measures plasma properties of coronal protons, ions, and electrons between 1.5 and 10 solar radii. Combines “occultation” with spectroscopy to reveal the solar wind acceleration region! slit field of view: Mirror motions select height UVCS “rolls” independently of spacecraft 2 UV channels: 1 white-light polarimetry channel LYA (120–135 nm) OVI (95–120 nm + 2nd ord.)

What produces “emission lines” in a spectrum? There are 2 general ways of producing extra photons at a specific wavelength. Both mechanisms depend on the quantum nature of atoms: “bound” electrons have discrete energies . . . The incoming particle can be either: Incoming particle Electron absorbs energy Energy re-emitted as light A free electron from some other ionized atom (“collisional excitation”) A photon at the right wavelength from the bright solar disk (“resonant scattering”) There is some spread in wavelength

Using lines as plasma diagnostics The Doppler effect shifts photon wavelengths depending on motions of atoms: If profiles are shifted up or down in wavelength (from the known “rest wavelength”), this indicates the bulk flow speed along the line-of-sight. The widths of the profiles tell us about random motions along the line-of-sight (i.e., temperature!) The total intensity (i.e., number of photons) tells us mainly about the density of atoms, but for resonant scattering there’s also another “hidden” Doppler effect that tells us about the flow speeds perpendicular to the line-of-sight.

UVCS results: over the poles (1996-1997 ) The fastest solar wind flow is expected to come from dim “coronal holes.” In June 1996, the first measurements of heavy ion (e.g., O+5) line emission in the extended corona revealed surprisingly wide line profiles . . . Off-limb profiles: T > 200 million K ! On-disk profiles: T = 1–3 million K

Solar minimum UVCS results UVCS has led to new views of the collisionless nature of solar wind acceleration. Key results include: The fast solar wind becomes supersonic much closer to the Sun (~2 Rs) than previously believed. In coronal holes, heavy ions (e.g., O+5) flow faster and are heated hundreds of times more strongly than protons and electrons, and have anisotropic temperatures. (Kohl et al. 1997, 1998, 2006)

Coronal holes: over the solar cycle Even though large coronal holes have similar outflow speeds at 1 AU (>600 km/s), their acceleration (in O+5) in the corona is different! (Miralles et al. 2001) Solar minimum: Solar maximum:

Waves: remote-sensing techniques The following techniques are direct… (UVCS ion heating was more indirect) Intensity modulations . . . Motion tracking in images . . . Doppler shifts . . . Doppler broadening . . . Radio sounding . . . Tomczyk et al. (2007)

Alfvén waves: from Sun to Earth Velocity amplitudes of fluctuations measured (mainly) perpendicular to the background magnetic field.

Visible-light coronagraph “blobs” SOHO/LASCO “wavelets” (Stenborg & Cobelli 2003)

Do blobs trace out the slow wind? The blobs are very low-contrast and thus may be passive “leaves in the wind.” Sheeley et al. (1997)

Polar plumes and jets Dense, thin flux tubes permeate polar coronal holes. They live for about a day, but can recur from the same footpoint over several solar rotations. Short-lived “polar jets” are energetic events that appear to eject plasma into the solar wind. Hinode/XRT (DeForest et al. 1997)

Streamers with UVCS Streamers viewed “edge-on” look different in H0 and O+5 Ion abundance depletion in “core” due to grav. settling? Brightest “legs” show negligible outflow, but abundances consistent with in situ slow wind. Higher latitudes and upper “stalk” show definite flows (Strachan et al. 2002). Stalk also has preferential ion heating & anisotropy, like coronal holes! (Frazin et al. 2003)

Coronal mass ejections Coronal mass ejections (CMEs) are magnetically driven eruptions from the Sun that carry energetic, twisted strands of plasma into the solar system . . . solar flare prominence eruption

The solar chromosphere

The chromospheric network Large numbers of lines (Mg, Ca; also IR, mm continuum) constrain the temperature “plateau.” There is more heating in the network “lanes” than in the cell-center regions. Vernazza, Avrett, & Loeser (1981) A: darkest cell-center C: avg. quiet Sun F: brightest network lane Leighton (1963)

Strongest fields in supergranular “funnels?” Peter (2001) Fisk (2005) Tu et al. (2005)

The solar photosphere

Photosphere: the “lower boundary” Photosphere displays convective motion on a broad range of time/space scales: β << 1 β ~ 1 β > 1

Summary Heating is everywhere . . . . . . and everything is in motion

Extra slides . . .

First observations of “stellar outflows” Coronae & Aurorae seen since antiquity . . . “New stars” 1572: Tycho’s supernova 1600: P Cygni outburst (“Revenante of the Swan”) 1604: Kepler’s supernova in “Serepentarius”

Open flux tubes: global model Photospheric flux tubes are shaken by an observed spectrum of horizontal motions. Alfvén waves propagate along the field, and partly reflect back down (non-WKB). Nonlinear couplings allow a (mainly perpendicular) cascade, terminated by damping. (Heinemann & Olbert 1980; Hollweg 1981, 1986; Velli 1993; Matthaeus et al. 1999; Dmitruk et al. 2001, 2002; Cranmer & van Ballegooijen 2003, 2005; Verdini et al. 2005; Oughton et al. 2006; many others!)

Doppler dimming & pumping After H I Lyman alpha, the O VI 1032, 1037 doublet are the next brightest lines in the extended corona. The isolated 1032 line Doppler dims like Lyman alpha. The 1037 line is “Doppler pumped” by neighboring C II line photons when O5+ outflow speed passes 175 and 370 km/s. The ratio R of 1032 to 1037 intensity depends on both the bulk outflow speed (of O5+ ions) and their parallel temperature. . . The line widths constrain perpendicular temperature to be > 100 million K. R < 1 implies anisotropy!

The Need for Better Observations Even though UVCS/SOHO has made significant advances, We still do not understand the physical processes that heat and accelerate the entire plasma (protons, electrons, heavy ions), There is still controversy about whether the fast solar wind occurs primarily in dense polar plumes or in low-density inter-plume plasma, We still do not know how and where the various components of the variable slow solar wind are produced (e.g., “blobs”). (Our understanding of ion cyclotron resonance is based essentially on just one ion!) UVCS has shown that answering these questions is possible, but cannot make the required observations.

Future diagnostics: more spectral lines! How/where do plasma fluctuations drive the preferential ion heating and acceleration, and how are the fluctuations produced and damped? Observing emission lines of additional ions (i.e., more charge & mass combinations) would constrain the specific kinds of waves and the specific collisionless damping modes. Comparison of predictions of UV line widths for ion cyclotron heating in 2 extreme limits (which UVCS observations [black circles] cannot distinguish). Cranmer (2002), astro-ph/0209301

Future Diagnostics: electron VDF Simulated H I Lyman alpha broadening from both H0 motions (yellow) and electron Thomson scattering (green). Both proton and electron temperatures can be measured.

Future Diagnostics: suprathermal tails Measuring non-Maxwellian velocity distributions of electrons and positive ions would allow us to test specific models of, e.g., velocity filtration, cyclotron resonance, and MHD turbulence. (Also these “seed particles” allow us to test models of SEP acceleration…) Cranmer (1998, 2001)