Lecture 5 The Formation and Evolution of CIRS

Transcription

1 Lecture 5 The Formation and Evolution of CIRS

2 Fast and Slow Solar Wind Fast solar wind (>600 km/s) is known to come from large coronal holes which have open magnetic field structure. The origin of slow solar wind (<400 km/s) is less certain. It is related to regions near magnetically closed coronal structures and the streamer belt. Schematic of locations sources of fast and slow solar wind during the declining and minimum phases of the solar cycle.

3 Parameters of the Solar Wind 1 3 E k = npmpv p where m p is the proton mass, v p is the 2 bulk velocity, n p is the density. E g = npv pgmpm s RS where M s is mass of the Sun, R S is the solar radius, G is the gravitational constant. T p and T e are the proton and electron temperatures and He 2+ /H + is the abundance ratio

4 Ulysses Observations of the Solar Wind Velocity Plasma measurements taken between 1992 and [McComas et al., 1998]. The velocity increases from about 450 km/s at the equator to about 750 km/s above the poles. Above 50 0 only fast solar winds streaming out of coronal holes were observed. Up to about 30 0 a recurrent CIR was observed with a period of about 26 days.

5 High Speed Streams The flow speed varies from pre-stream levels (400 km/s) reaching a maximum value (600 km/s 700 km/s) in about one day. The density rises to high values (>50 cm -3 ) near the leading edges of the streams and these high densities generally persist for about a day. The peaks are followed by low densities lasting several days. The proton temperature varies like the flow speed. The high speed streams tend to have a dominant magnetic polarity. The dominant source of high speed streams is thought to be field lines that are open to interplanetary space. These regions are known as coronal holes.

6 Solar Wind Structure Solar Minimum and Maximum Solar minimum fast tenuous flows from coronal holes at high and middle latitudes, denser more variable flows near the equator. Solar maximum slow to intermediate flows, CMEs, no large holes but high speed in small as well as large holes.

7 Solar Cycle Dependence of HCS The waviness of the current sheet increases at solar maximum. The current sheet is rather flat during solar minimum but extends to high latitudes during solar maximum. During solar minimum CIRs are confined to the equatorial region but cover a wide range of latitudes during solar maximum. The average velocity of the solar wind is greater during solar minimum because high-speed streams are observed more frequently and for longer times

8 High Speed Streams Observed in the Inner Heliosphere (0.29AU to 1AU) Observations near solar minimum from Helios in ecliptic (Balogh et al., 1999). Helios passed through the HCS (B φ plot) as the Sun rotated. Slow solar wind in HCS (dashed lines). High speed streams have lower density. No CIR-associated shocks.

9 Organizing the Solar Wind Velocity Solar wind velocity for 2004 as a stack plot 30 day segments advanced 27 days in successive rows (McPherron and Weygand, 2006). Triangles indicate stream interfaces (SI). A persistent SI occurred on about the 15 th of each rotation.

10 Coronal Structure Near Solar Minimum The coronal magnetic field is thought to be important for understanding the structure of the corona. No measurements of the coronal field structure. Only line of sight component. Use measured photospheric field to infer coronal field. Include photospheric field in MHD simulations to create models of the solar corona. (B r on R s from synoptic magnetic field observations NSOKP and WSO or full-disk magnetograms.) (Mikić et al., 1999). Outer domain was at r=30r S. Initially a quasi-steady B r = ϕ coronal model was used for a given B r a model field is calculated in the coronal A spherically symmetric wind solution was used to specify pressure, mass density and velocity on the surface. Integrate to steady state. The Lundquist number was (10 13 in quiet corona). Test simulation results against white light observations of coronal field.

11 Comparison of simulations and observed polarization brightness Use density from MHD and scattering function to calculate brightness. Agreement is good. Streamer belt is on closed field lines.

12 Simulations and Observations of Polarization Brightness

13 Images of Sun in EUV and X-Rays Coronal hole appears as a black region cooler less dense than rest of solar disk. Calculate open field line regions (black) and closed field line regions (gray). Compare with SOHO EIT (EUV) observations. Observed coronal hole extends from the pole to the equator Observed coronal hole is region of open field lines. Photospheric magnetic field seems to control streamer belt and coronal holes.

14 Comparing MHD Simulations with Disk Images

15 From the Solar Wind Back to the Sun If the model approximates the true field well, it can be used to investigate source locations of solar wind features. Mapping of solar wind into MHD domain ballistically and then used MHD to map source to Sun. Used Wind and Ulysses observations. Slow solar wind to boundaries of hole. Fast solar wind to center.

16 Time Dependence in the Corona and Solar Wind Photospheric magnetic field was changed at 10 times real speed can t study detailed changes. Coronal hole maps black open, gray closed field line regions (left). Field line traces blue into Sun, red out (second left). Polarization brightness (third left). Shape of current sheet (right). Closed field regions sporadically open. Streamer belt plasma is carried out into the solar wind. Periodic reconnection at coronal hole boundaries?

17 Evolution over 15 Carrington Rotations

18 Evolution over 15 Carrington Rotations

19 Observations of High Speed Streams Velocity, density and proton temperature of two high speed streams Speed and temperature have similar variations with time Note that low speed corresponds to high density and vice versa Flow Speed (km/s) Density (cm -3 ) Temp. (K)

20 The Archimedean spiral associated with slow streams is curved more strongly than for a fast stream. For B ϕ B r 2 r0 0 = B r vϕ rω = v r Sun Br rω > Sun v ϕ Bϕ rω SunBr vr Forming a CIR tanψ = B ϕ B r Because field lines are not allowed to intersect at some point an interaction region develops between fast and slow streams. Since both rotate with the Sun these are called corotating interaction regions (CIR). On the Sun there is an abrupt change in the solar wind speeds but in space the streams are spread out.

21 At the interface between fast and slow streams the plasma is compressed. The characteristic propagation speeds (the Alfven speed and the sound speed) decrease. At some distance between 2AU and 3AU the density gradient on both sides of the CIR becomes large and a pair of shocks develop. The shock pair propagate away from the interface. The shock propagating into the slow speed stream is called a forward shock. The shock propagating into the fast wind is called a reverse shock Forming a CIR

22 Time series of parameters associated with a CIR Between the two shock waves, and centered on the interface, the plasma is compressed This implies a higher density of S plasma than unshocked S plasma Similarly the shocked F plasma is higher density than unshocked F plasma, but the density of F < density S since fast plasma has lower density than slow plasma The S plasma is moving faster than S, but slower than F which is slower still than F The S plasma has a positive azimuthal velocity, the interface a zero azimuthal velocity, and the F a negative azimuthal velocity The magnitude of the magnetic field is compressed between the shocks There is increased magnetic turbulence and temperature in the interaction region Not shown is a tipping of the IMF out of the ecliptic plane

23 A Typical CIR at Earth Orbit CIRs with well developed shocks are very rare at Earth orbit. The dashed line shows the interface between the high and low speed streams. Dst is only about - 70nT.

24 Corotating Interaction Regions at Jupiter (5AU) (Joy et al., 2002) In the outer heliosphere the CIRs can steeped and have shocks.