P472G. The Solar Wind

Transcription

1 P472G The Solar Wind

2 Copyright No9ce The material presented during this course may contain items collected from third- party sources and are presented to you for your personal study. The US Copyright Act of 1976 governs reproduc9on of copyrighted material. Parts of this presenta9on may contain material that is protected under this act. This material has been reproduced for you in accordance with the Act s fair use provision that allows educa9onal ins9tu9ons to furnish por9ons of copyrighted material to students. Materials reproduced under this provision may not be used for any purpose other than private study, scholarship or research. Electronic copies should not be shared with unauthorized third par9es.

3 History: Geomagne9c Storms Tsurutani et al. JGR 1995 Magne9c Storms

4 Correla9on Between Magne9c Storms and Solar Ac9vity For over one hundred years people no9ced that geomagne9c storms tended to occur with an approximate 27 day period. This plot shows a study published in 1962 showing a correla9on between sunspot numbers and recurrent ac9vity (M). Note that the M regions seem to peak during the diminishing por9on of the sunspot cycle. Some geomagne9c storms do not seem to be associated with the solar cycle, (S). These are probably due to flares that did not persist for a full solar rota9on.

5 Observa9ons of Comets in Interplanetary Space As a comet approaches the sun the heat of the sun begins to cause the surface to break apart and gases escape from the comet. While the finite size dust par9cles con9nue to move approximately along the cometary orbit, the gas tail points away from the sun at all por9ons of the orbit. Biermann (1934) argued that the gas tail showed evidence that the sun emi`ed a steady stream of par9cles that would cause the gas trail to be directed away from the sun.

6 Comet Lovejoy NASA/ESA SOHO solar observatory, 17 December 2011

7 Parker s Model (1958)

8 Parker s Solu9ons (Warwick Dept of Physics, PX420) Type I and II are double valued and thus, non- physical Type III has supersonic speeds at the solar surface. These are not observed. Type IV: a solar breeze, a subsonic flow Type V: a unique solu9on; begins subsonic and passes through Cs at r c for C=3 to become supersonic. This is the solar wind solu9on. Rc = GMo/2C 2 ~ 9 to 10 Ro.

9 Parker s Solu9ons

10 The First Solar Wind Measurements While the first satellite measurements of the solar wind were made with the Russian Soyuz spacecraft, the first continuous measurements were first made by an electrostatic analyzer and a magnetometer onboard Mariner II during its epic 3-month journey to Venus in 1962 provided firm confirmation of a continuous solar wind flow and spiral heliospheric magnetic field that agree with Parker s model, on average. Mariner II also showed that the solar wind was highly variable, being structured into alternating streams of high and low-speed flows that lasted for several days each. The observed magnetic field was also highly variable in both strength and orientation.

11 Radial Profiles of Solar Wind Speed and Density Pioneer 10 Spacecran

12 Evolu9on of the Solar Wind Stream Structure with Distance Spa9al variability of the solar wind ouolow and solar rota9on produce radial varia9ons in speed. Faster wind overtakes slow wind ahead while outrunning slow wind behind. As a result, the leading edges of high- speed streams steepen with increasing heliocentric distance. Plasma is compressed on the leading edge of a stream and rarefied on the trailing edge. The major accelera9ons and decelera9ons of the wind then occur at the shocks and the stream profile becomes a damped double sawtooth. Because the sound speed decreases with increasing heliocentric distance, virtually all high- speed streams eventually have shock pairs on their leading edges.

13 Solar Wind Magne9c Field A parcel of coronal material has a radial velocity of u(r) in the r direc9on when observed from either the rota9ng or the non- rota9ng frame. But in the rota9ng frame it acquires an `apparent tangen9al velocity component v (r) in the Φ direc9on, given by v (r) = - ωr Where ω is the Sun s angular velocity of rota9on. Now because of the frozen- in condi9on, the magne9c field lines in the solar wind must be parallel to the plasma flow velocity (in the rota9ng frame), so that in the rota9ng frame B Φ /B r = v/u = - ωr/u(r) We are trea9ng this as a 2D problem (in a plane through the solar equator). The differen9al equa9on for the field lines in the 2D rota9ng frame is obtained from - ωr/u(r) = v/u = (rdφ/dt)/(dr/dt) So rdφ/dr = - ωr/u(r)

14 Solar Wind Magne9c Field In regions where the solar wind speed is essen9ally constant with r (far away from the sun) then dφ/dr = - ω/u or, inver9ng, dr/dφ = - ωr/u(r) and, integra9ng, r = - (u/ω)φ + K Evalua9ng K at the solar surface, K = R + (u/ω)φ o So the field line sa9sfies the equa9on (r- R) = (u/ω)(φ o - Φ)

15 Solar Wind Magne9c Field Since the magne9c field lines are carried out by the plasma flow, the conserva9on of magne9c flux requires that B r = B o (R/r) 2. Then the tangen9al component of B is B Φ = - B o (ωr/u)(r/r) 2 That is, the equatorial field is in the form of a spiral. Tanθ = v/u = - ωr/u At 1 AU, θ ~ 45 o

16 The Solar Wind at 1 AU Since ~1960 satellites have observed values for some key parameters of the solar wind observed at a radial distance of 1 AU. However, it must be emphasized the solar wind is highly variable, so that the concept of a representa9ve `average value is to be taken with care: Proton Density 6.6cm- 3 Electron Density 7.1cm- 3 He2+ Density 0.25cm- 3 Flow Speed 450km/s Proton Temperature 1.2x10 5 K Electron Temperature K

17 The Solar Wind at 1 AU These plots give another illustra9on of the characteris9cs of the solar wind. In this case they are probability density func9ons of speed, electron number density, and log- temperature. The y- axes are all presented using a log scale. The probability density func9ons are scaled so that the total area under them one. This just means that there is a probability of one that the speed in some region will have some value.

18 The Solar Wind at 1 AU: Velocity 60 days of solar wind velocity data for a period near solar maximum. 60 days of solar wind velocity data near solar minimum. Note that the scales differ slightly.

19 The Solar Wind at 1 AU: IMF The interplanetary magne9c field (IMF) for 60 days near solar maximum. The interplanetary magne9c field (IMF) for 60 days near solar minimum. Note, the scales are much different with the largest values occurring near solar maximum.

20 The Solar Wind at 1 AU: Temperature The solar wind temperature for 60 days near solar maximum. The solar wind temperature for 60 days near solar maximum. The differences are not drama9c.

21 Solar Wind (1 AU) and Corona Temperatures and Electron Densi9es

22 The Solar Dipole Field The radial stretching of an otherwise large- scale dipolar magne9c field is shown here. The most drama9c effect is in the equatorial plane, where the radial expansion of the solar wind blows straight across an undistorted dipolar field. This configura9on would be typical of the corona at solar minimum. However, the solar wind has the same effect on any coronal loops.even the numerous small ones that appear at solar maximum.

23 Ulysses Space Cran Solar Orbit The joint ESA/NASA `Ulysses spacecran performed 3 orbits of the Sun between 1992 and A Jupiter flyby was used to achieve an orbital plane perpendicular to the eclip9c, that carried the spacecran over both solar poles. These are the only direct observa9ons that we have on how the solar wind depends on solar la9tude.

24 Ulysses Solar Satellite Solar Minimum Condi9ons Ulysses first orbit was during solar minimum. A very clear picture emerged: fast and consistent solar wind speed was seen at all la9tudes except the equatorial zone, where the slow wind prevailed. The solar magne9c field was strongly dipolar: outward at northern la9tudes, inward at southern la9tudes. However the r2- scaled density increased in equatorial la9tudes, so that the solar wind dynamic pressure remained essen9ally constant.

25 Ulysses Solar Satellite Solar Maximum Condi9ons Ulysses second orbit was during solar maximum. In this case a much more confused arrangement of fast and slow solar wind was observed. This is not surprising; we know the coronal structure is much more complex at solar maximum.

26 Wobbling Current Sheet This figure shows a smooth dipolar solar field stretched flat into a heliocentric current sheet by the solar wind. The axis of the current sheet is 9lted rela9ve to the sun s rota9onal axis and so `wobbles as the sun rotates. In this configura9on the Earth would spend half its 9me above the sheet and half below it.

27 Wobbling Current Sheet Even with the wobble, Earth is always near the heliospheric current sheet - so IMF field lines near Earth are typically oriented in a direc9on that is either toward or away from the Sun, plus a 45 rota9on due to the Parker spiral. For a simple dipole field, Earth would cross the current sheet once per solar rota9on. spending half its 9me in a `sector where the field lines point toward the Sun, and half in a sector where the line point away from the Sun.

28 Overall Shape of the Interplanetary Magne9c Field The wobble mo9on due to the 9lt of the Sun s dipole axis together with the possibly non- dipolar magne9c topology at the photosphere act together to distort the heliospheric current sheet away from being a simple flat plane. Instead, it is onen described as being like a wavy `ballerina s skirt, spinning with the Sun s rota9on. Here we see a wave introduced by the Sun s north polar coronal hole extending to northern mid- la9tudes.

29 The Solar IMF

30 Next Time: Solar Wind Interac9ons with the Planets Next 9me: Solar wind interac9ons with the planets. Magne9zed planets have `magnetospheres Unmagne9zed planets have solar wind `shadows