The solar wind (in 90 minutes) Mathew Owens

The solar wind (in 90 minutes) Mathew Owens 5 th Sept 2013 STFC Advanced Summer School m.j.owens@reading.ac.uk

Overview There s simply too much to cover in 90 minutes Hope to touch on: Formation of the solar wind and the Heliosphere Parker Spiral Three-dimensional solar wind Structure: Corotating interaction regions and transients Solar wind variability over centuries/millenia Apologies to: Turbulence and fine-scale structure Space weather Magnetospheric and ionospheric physics Solar-climate studies 2

Sunspots evidence of photospheric structure

Eclipses evidence of a solar atmosphere STFC Summer School 2009 4

Carrington - evidence for particles from the Sun 11:18am, 1 September 1859, Redhill Richard Carrington: observed a very bright flare on the Sun Next day: auroras down to very low latitudes (Cuba, Hawaii) Telegraph systems disrupted Evidence for particles from the Sun, travelling at ~2000 km/s to the Earth Corpuscular radiation STFC Summer School 2009 5

Evidence of a constant wind - Comet tails evidence of a solar wind comets Biermann (1951): non-radial comet tails imply constant wind of few hundred km/s STFC Summer School 2009 6

Temperature (K) Coronal temperatures 10 6 3 10 6 K 1.5 10 6 K 1 10 6 K 10 5 6 10 4 K 10 4 0 2000 4000 6000 8000 Height above photosphere (km) 6 10 3 K braiding of magnetic field by photospheric motions is central to coronal heating

Origin of the solar wind Magnetic carpet of loops Most closed, some extend far above solar surface Plasma flows out along field lines into space STFC Summer School 2009 8

Solar wind source: corona Hot corona has much higher pressure than interstellar medium Thermal scale height greater than gravitational scale height Can t balance pressures with a static corona (Parker, 1957) Corona expands into space, forming the solar wind Note: in lower corona, plasma <<1: plasma follows field STFC Summer School 2009 9

Parker s solution Parker(1958) was the first to propose a model of the solar wind assuming a steady flow of plasma independent of time, as opposed to a static corona. He began from the mass and momentum conservation equations, taking time derivatives as zero since considering a steady flow. u 0 u. u p j B F g Found a solution of the form u 2 2kBT 1 du 4 m u dr kbt mr GM r 2 S

Possible solutions of the Parker solar wind equation

First solar wind observations Mariner 2 in 1962 12

The heliosphere The cavity in our local interstellar wind termination shock heliopause interstellar wind dominated by the solar wind and Sun s magnetic field heliosheath bow shock

MHD simulations of the heliosphere Solar wind is confined in a cavity in interstellar space called the heliosphere that surrounds all the planets of the solar system Density is enhanced behind bow shock in heliosheath Termination shock heliopause heliosheath interstellar wind bow shock

A stellarsphere Hubble observations of the heliosheath behind the bow shock where the heliosphere of LL Ori heliosphere meets its (dense) local interstellar wind in the Orion nebula heliosheath bow shock interstellar wind heliopause

A stellarsphere Gemini adaptive optics observations near the galactic centre bow shock fast-moving star galactic centre heliosheath

Frozen-in theorem: The convective limit THE INDUCTION EQUATION B/ t = (V B ) + 2 B /( o ) convective term MAGNETIC REYNOLDS NUMBER diffusive term R m = { (V B ) } / { 2 B /( o )} { (V B ) } V c B c / L c ; { 2 B /( o ) } {B c / L c 2 } {1/( o )} Thus R m o V c L c Region (mhos m -1 ) V c (m s -1 ) L c (m) R m Base of corona 10 3 10 5 10 6 10 8 S.W. @ 1AU 10 4 10 5 10 9 10 12 Thus R m o V c L c >> 1 so B/ t = (V B ) this is the convective limit & leads to the frozen-in flux theorem

Frozen-in theorem: The convective limit ( by definition of B ) Charged particle motions Magnetic Field B Lorentz Force B B

Parker spiral Solar wind flow is radial Sun Solar rotation and radial solar wind generates a Parker spiral field structure 1 Field is frozen-in to the solar wind flow

Interplanetary scintillation

Interplanetary scintillation Solar rotation and radial solar wind generates a Parker spiral field structure Co-rotates with the solar corona (every 27 days in Earth s frame)

X-ray observations of flare particles Tracked by Ulysses satellite, looking down on ecliptic plane from over solar pole, using X-rays generated by super-thermal flare electrons moving along the field lines through the thermal solar wind

Calculating the Parker Spiral t left Sun at time t V SW t left Sun at time t+ t r r r t tan = r t / (V SW t) = tan -1 (r /V SW ) Sun

Observations of the spiral angle Y X X V sw = 332-358 km s -1 Y X X V sw = 379-403 km s -1 Colour histograms: Observed distributions of for 1965-2002, divided into 9 V SW ranges giving equal numbers of samples (grey) range and (black) average predicted by frozen-in = tan -1 {V SW /(r )} Y V sw = 430-463 km s -1 Y V sw = 507-580 km s -1 Spiral unwinds for higher V sw - as predicted by thory

Observations of the spiral angle V sw = 332 358 km s -1

Parker spiral angle at 1 AU 39

Solar wind effect on the corona The heliospheric magnetic field is a result of the Sun s magnetic field being carried outward, frozen in to the solar wind. Within the corona, the magnetic field forces dominate the plasma forces. As the field strength decreases with distance, beyond the Alfvén radius at a few solar radii, the plasma flow becomes dominant, and the field lines are constrained to move with the solar wind. Model of Pneumann and Kopp (1971)

Modelling the corona 41

The coronal source surface 42

PFSS solutions Magnetic field polarity at coronal source surface 43

Ulysses

Three-dimensional structure of interplanetary magnetic field 45

Heliospheric Ulysses passages above the current maximum latitude sheet of the current sheet STFC Advanced Summer School m.j.owens@reading.ac.uk

Ulysses fast latitude scans 1 st - near sunspot min 2 nd - near sunspot max

B R invariance with latitude Radial flow at r > 10R S with B r independent of latitude V SW N Sun S Slight non-radial flow at r > 10R S to equalise P t & thus B r B r tangential pressure, P t B r2 / 2 o (as << 1)

Open Solar Flux Flux threading the coronal source surface closed field line Unsigned Flux, F U = + /2 2 - /2 0 B R r 2 cos( ) d d r = heliocentric distance B R = radial field = solar latitude = solar longitude open field lines STFC Advanced Summer School m.j.owens@reading.ac.uk

Open solar flux 50

Solar wind speed Sunspot minimum FAST Solar Wind in coronal holes SLOW Solar Wind in equatorial streamer belt Red = inward field Blue = outward field

Close to solar minimum the flow pattern close to the Sun can be approximated as a band of slow wind at low latitudes, centred on the Sun s dipole equator, with fast wind at all higher latitudes. This pattern of fast and slow solar wind is occasionally disturbed by transient flows associated with coronal mass ejections. Pizzo (1991) STFC Advanced Summer School m.j.owens@reading.ac.uk

Solar cycle evolution of speed

Characteristics of slow and fast solar wind Property at 1 AU Slow wind Fast wind Speed (v) ~400 km/s ~750 km/s Number density (n) ~10 cm 3 ~3 cm 3 Flux (nv) ~3 10 8 cm 2 s 1 ~2 10 8 cm 2 s 1 Magnetic field (Br) ~3 nt ~3 nt Proton temperature (Tp) ~4 10 4 K ~2 10 5 K Electron temperature(te) ~1.3 10 5 K (>Tp) ~1 10 5 K (<Tp) Composition (He/H) ~1 30% ~5% STFC Advanced Summer School m.j.owens@reading.ac.uk

Ionisation states Heavy ions (oxygen, carbon, magnesium) occur in small quantities in solar wind Ionisation states determined by temperature of plasma when collisional, in corona When collisionless, ionisation stares don t change Freezing-in temperature temperature of solar wind when collisions stop Diagnostic of conditions in corona where solar wind originates Very different in fast and slow solar wind STFC Summer School 2009 55

Loop opening

Imaging the solar wind 57

Interaction of fast and slow wind If fast and slow solar wind streams at same latitude, fast can overtake slow Interaction: compression, can drive shocks 58

V (km/s) V T, V N (km/s) n p (cm -3 ) T p (x10 5 K) B (nt) B, B P Total (npa) 800 700 600 500 400 50 0-50 -100-150 2.0 1.5 1.0 0.5 0.0 4 3 2 1 0 3 2 1 0 360 270 180 90 0-90 10-2 10-3 10-4 10-5 V T V N FS HCS SI RW 332 333 334 335 336 337 338 Day of 1992 B B In situ observations 59

Transient events - CMEs 60

In situ observations 61

Interaction with magnetosphere 62

Generation of energetic particles

Suprathermal electrons Owens and Crooker, 2008

CMEs add magnetic flux to the heliosphere 65

Cosmic rays Sunspot Number Huancauyo Hawaii neutron monitor counts (>13GV) Climax neutron monitor counts (>3GV)

14 C & 10 Be: spallation products from O, N & Ar 14 C 1/2 = 5370 yr < q G > = 2 atoms cm -2 s -1 GALACTIC COSMIC RAYS STRATOSPHERE ( 2/3) 10 Be 1/2 = 1.5 10 6 yr < q G > = 0.018 atoms cm -2 s -1 14 C+0 14 C0 ; 14 C0+0H 14 C0 2 + H TROPOSPHERE ( 1/3) OCEANS ( ~1 year) ( ~1 week) 10 Be + AEROSOL BIOMASS ICE SHEETS

HMF reconstructions 75

Solar Modulation Parameter, (MV) Millennial Variation composite (40-year means) from cosmogenic isotopes 14 C & 10 Be 1000 we are still within recent grand maximum 800 600 400 200 0-6000 -4000-2000 0 2000 Year AD

Open questions Where does slow solar wind come from? How is the corona heated and the solar wind accelerated? What is the structure of the heliopause? How do CMEs mediate the solar cycle? How does the Sun evolve over centennial and millennial timescales? 77

The Future Solar Orbiter ESA-led mission Close to the Sun (0.23 AU) Orbit inclination to ecliptic (30 degrees) Launch in 2017? Lots of UK involvement (magnetometer at Imperial,, electron detector at MSSL) 78