行星際激波在日球層中的傳播：Propagation of Interplanetary Shocks in the Heliosphere （第二部分）

行星際激波在日球層中的傳播：Propagation of Interplanetary Shocks in the Heliosphere （第一部分）- Chapter 1: Introduction & Chapter 2: Basics of Magnetohydrodynamics

行星際激波在日球層中的傳播：Propagation of Interplanetary Shocks in the Heliosphere （第三部分）- Chapter 4: Analysis methods

行星際激波在日球層中的傳播：Propagation of Interplanetary Shocks in the Heliosphere （第四部分）- Chapter 5: Discussion & Chapter 6: Summary and Conclusions

Propagation of Interplanetary Shocks in the Heliosphere?
[ Chapter 3 ]

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行星際激波在日球層中的傳播：Propagation of Interplanetary Shocks in the Heliosphere （參考文獻部分）-CSDN博客

? ?Chapter 3? ?Solar and heliospheric physics? ?

The Sun?? is the ??main source of energy for all lives on the earth?? as well as the ??main defining object of the solar system dynamics??. Like all other stars, the Sun has both ??inner complex dynamics?? and ??influence on interplanetary space??. In this chapter, for the sake of completeness and to explain the origin of these ??tremendous dynamics??, the ??inner structure of the Sun?? as well as some important ??heliospheric activities?? are briefly discussed.

太陽??作為??地球上所有生命的主要能量來源??，同時也是??太陽系動力學的主要決定者??。與其他恒星類似，太陽同時具有??復雜的內部動力學過程??和??對行星際空間的深遠影響??。本章將系統闡述太陽的??內部結構??及其??重要的日球層活動??，以完整揭示這些??驚人動力學現象??的起源機制。

? ▍3.1? The Sun?▍

Our star, the Sun, is a ??G2-V type star?? (Aschwanden, 2014). All stars are in a ??balance of opposing forces?? between ??outward radiation fueled by nuclear fusion?? and ??inward pressure dictated by gravitational force??. The ??radiative force?? is the product of the ??fusion of hydrogen atoms, which constitutes about 73 %, into helium atoms, which constitute 25 % of the Sun?? (Basu and Antia, 2007). This process is a ??thermonuclear reaction?? via turning a ??proton into a neutron??.

Structurally, the Sun has ??inner zones?? and ??outer zones??. The former consists of the ??core, radiative zone, interface layer (tachocline), and convection zone??. The latter consists of the ??photosphere, the chromosphere, the transition region, and the corona?? (NASA, 2017) as shown in ??Figure 3.1??.

太陽屬于??G2-V型主序星??（Aschwanden，2014）。所有恒星都維持著??兩種對立力量的平衡??：??由核聚變驅動的向外輻射壓力??與??引力作用導致的向內收縮壓力??。其中，??輻射壓力??源自??氫原子（占73%）聚變為氦原子（占25%）??的過程（Basu與Antia，2007），這一??熱核反應??本質上是通過??質子轉化為中子??實現的。

從結構上看，太陽可分為??內部區域??和??外部大氣??。??

內部區域??包含：核心區??、??輻射層??、??差旋層（tachocline）??、??對流層?。

?外部大氣??則由以下層次構成：??光球層??、??色球層??、??過渡區??、??日冕層??（NASA，2017），具體結構如??圖3.1??所示。

Figure 3.1: The layers of the Sun (Credit: NASA, 2017)

The ??core of the Sun

The ??core of the Sun?? is ??25% of the solar radius?? (Garc′?a et al., 2007), the ??density is about 150g/cm3?? (Basu et al., 2009) contains ??34% of the solar mass?? even though only ??0.8% of the total volume??, and it is the ??hottest part of the Sun?? that a ??temperature of 15 million Kelvins?? (Dr. David H. Hathaway, 2015). This is the zone where the ??thermonuclear reaction is going??. For the ??radiative zone??, it starts from the ??edge of the core to the interface layer (tachocline)??, and much of the ??energy generated in the core?? is carried away by ??photons?? though photons take ??a million years?? to reach the next layer due to the ??dense material of the region?? (Chemin, 2022).

??太陽核心??占據??太陽半徑的25%??（García等，2007），其??密度高達150g/cm3??（Basu等，2009），雖然僅占??太陽總體積的0.8%??，卻包含??34%的太陽質量??。作為太陽??最熾熱的區域??，核心溫度達到??1500萬開爾文??（Hathaway，2015），是??持續進行熱核反應的區域??。在??輻射區??（從核心邊緣延伸至差旋層），??核心產生的能量??主要通過??光子傳遞??，但由于該區域??物質密度極高??，光子需??耗費百萬年??才能抵達下一層（Chemin，2022）。

Interface Layer

The ??very next layer, the interface layer (tachocline)?? is ??much of our interest?? because it is believed that the ??solar magnetic dynamo originates in this very thin layer?? (Tobias and Weiss, 2007) that ??magnetic field lines can become stronger?? due to the ??changes in the fluid flow velocities crossing this layer??. However, recent ??radio observations of brown dwarf stars?? show that, despite ??not having a tachocline layer??, they can have ??similar magnetic strength and activities just like the Sun??, which indicates the ??convection zone may be solely responsible for the solar dynamo?? (Route, 2016).

??差旋層（tachocline）??作為核心與對流層的過渡區具有特殊意義。這個??極薄的界面層??被認為是??太陽磁發電機效應的起源地??（Tobias與Weiss，2007），??因流體速度突變而增強的磁力線??在此形成。但最新研究顯示，??缺乏差旋層的褐矮星??同樣存在??類太陽磁活動??（Route，2016），暗示??對流層可能獨立維持磁發電機效應??。

??Convection Zone??

Next, the ??convection zone?? is the ??outermost layer of the inner zones?? and it extends about ??200,000 km from the depth?? and the ??temperature is approximately 2 million Kelvins?? (Christensen-Dalsgaard et al., 1991), which makes the zone ??relatively cooler for heavy ions to keep their electrons?? so that the zone is ??opaque to the radiation?? (Ortu?no-Araujo et al., 2012), which in turn ??traps the heat?? and eventually the ??fluid becomes unstable and start convecting??. As a result, this layer is ??very turbulent?? (Brummell et al., 1995), which in addition to the ??rotational motion creates electric currents and magnetic fields??, and the ??gas pressure is much more dominant than the magnetic pressure?? in this region, and because of this the ??magnetic field is dragged and twisted by the fluid??, which then ??propagates passing through the photosphere, chromosphere, the transition region up to the corona?? and creates a ??multitude of activities on the surface of the Sun called solar activity??.?

??對流層??作為??太陽內部最外層??，延伸約??20萬公里??，溫度約??200萬開爾文??（Christensen-Dalsgaard等，1991）。該區域溫度使得??重離子能束縛電子??，導致??輻射不透明??（Ortu?o-Araujo等，2012），進而??形成熱約束??。當??流體失穩開始對流??時，會產生??劇烈湍流??（Brummell等，1995）。這種??旋轉運動??不僅生成??電流與磁場??，還使得??氣體壓力遠超過磁壓力??，導致??磁場被流體拖曳扭曲??。這些磁結構??穿越光球層、色球層、過渡區直達日冕??，最終引發??各類太陽表面活動現象??。

Photosphere

In the ??photosphere??, visible ??darker areas called sunspots?? appear due to ??localized strong magnetic regions??, which ??inhibit some of the heat from reaching the surface?? and this makes the areas ??cooler than the other parts of the Sun?? (Babcock and Babcock, 1955). Because of the mentioned ??strong magnetic nature??, the ??magnetic field lines around the sunspot often twist and cross??, causing a ??tension of energy that bursts in an explosion called a solar flare?? (Chakraborty and Basak, 2022).

在光球層??中，由于??局部強磁場區域??的存在，會形成可見的斑——太陽黑子??（Babcock父子，1955）。這些磁場??阻礙熱量傳遞至表面??，使得黑子區域??溫度顯著低于周邊??。這種??強磁場的扭曲特性??導致??黑子周圍磁力線頻繁交織??，最終??引發能量劇烈釋放的爆發性事件——太陽耀斑?**?（Chakraborty與Basak，2022）。

Corona

The ??corona?? is the ??hot and ionized outermost layer of the Sun?? and it is ??millions of Kelvins greater than that of the surface of the Sun?? (Aschwanden, 2006). The ??solar magnetic field confines much of the coronal plasma??, but some of the plasma ??spread with a supersonic speed into interplanetary space??, which is ??solar wind??. The ??solar magnetic field forms so-called flux tubes??. These ??ropes are tensed by the differential rotation of the Sun??. Therefore, ??store plenty of energy??. If a ??certain configuration appears (X-shape)?? the ??magnetic structure reorganizes??. Hence, a ??huge amount of energy is released??. The name of this ??20 million K-degree flash is a solar flare??. It emits ??plenty of energized protons and other ions in the heliosphere?? and very often, but not always a ??huge amount of solar plasma ejects to the IP space??. The process is called ??coronal mass ejections (CME)??. However, ??CME could occur without a flare too??.

These phenomena are the ??main drivers of space weather?? and the ??conditions of the terrestrial cosmic environment?? (Facsk′o et al., 2022). The ??solar activity is cyclic??, which is called the ??solar cycle??, and it is ??approximately 11 years?? (Center/SDO, 2015), see ??Figure 3.2??. Each cycle the ??magnetic field of the Sun flips??, and the periods, in which, the ??sunspots are the greater in number are called solar maximum?? and the ??lower in number are called solar minimum??.

??日冕??作為太陽??最外層的高溫電離大氣??，其溫度??比光球層高出數百萬開爾文??（Aschwanden，2006）。雖然??太陽磁場束縛著大部分日冕等離子體??，但仍有部分等離子體??以超音速進入行星際空間??，形成??持續不斷的太陽風??。日冕中的??磁通量管結構??因??太陽較差自轉??而不斷拉伸，從而??儲存巨大磁能??。當出現??特定的X型磁重聯構型??時，??磁場結構會發生重組??，瞬間釋放出??相當于2000萬開爾文能量的閃光——即太陽耀斑??，并向日球層噴射??大量高能粒子??。這類事件通常（但非必然）會伴隨??巨量等離子體拋射進入行星際空間??，該過程稱為??日冕物質拋射（CME）??。值得注意的是，??CME也可能獨立于耀斑單獨發生??（Facskó等，2022）。

這些爆發活動是??空間天氣的主要驅動源??，直接影響??地球空間環境狀態??。太陽活動呈現??約11年的周期性變化??（太陽活動周），其??全球磁場極性每周期發生一次反轉??（SDO研究中心，2015），參見??圖3.2??。??黑子數最多的時期稱為太陽活動極大期??，??最少的時期則稱為極小期??。

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Figure 3.2: The sunspot cycle over the last several decades (Figure is from Pasachoff et al., 2014; Figure 8)

?▍3.2 The Solar Wind and the Interplanetary Magnetic Field?▍

The Sun ejects highly energized and ionized charged particles continuously in all directions, which is the solar wind.?? The ??solar wind is a collisionless plasma?? that consists of ??equal amounts of protons and electrons?? with the addition of ??negligible 3 to 6% helium?? (Neugebauer, 1981). As a result, the ??solar wind is a quasi-neutral plasma?? whose characteristics can be defined by ??magnetohydrodynamics (MHD)??. In MHD, ??Alfvén theorem states?? that in a fluid with ??high electric conductivity?? the ??magnetic field line is frozen in it?? and moves along with it (Alfvén, 1942). Since the ??solar wind is one such fluid?? the ??Sun's magnetic field lines are frozen in the flow?? that they are ??forced to propagate with the solar wind?? (Roberts, 2007), forming the ??interplanetary magnetic field (IMF)??.

Depending on the ??origination point?? whether in the ??coronal holes?? or the ??equatorial belt of the Sun??, the solar wind is classified as the ??fast solar wind with a velocity of 750 km/s?? and the ??slow solar wind with a velocity of 400 km/s?? respectively (Feldman et al., 2005). The ??coronal holes are areas that appear dark in X-ray images?? because these areas are ??much cooler and less dense than other regions?? (Cranmer, 2009) and the ??magnetic field around these areas does not loop back down but extends into the interplanetary space?? (Parker, 1959), so the ??plasma can easily flow out, creating the fast solar wind??. The ?? holes can appear anywhere on the coronal areas during the solar maximum?? while they usually appear on ??northern or southern poles of the sun during the solar minimum?? (McComas et al., 2003).

3.2 太陽風與行星際磁場??

??太陽持續向各個方向噴射高能電離帶電粒子??，這就是????太陽風。??太陽風是一種無碰撞等離子體??，由??等量的質子與電子??組成，并含有??少量（3%-6%）的氦??（Neugebauer, 1981）。因此，??太陽風是準中性的等離子體??，其特性可以用??磁流體力學（MHD）??來描述。根據MHD中的??阿爾芬定理??，在??高電導率的流體??中，??磁力線會被凍結??并隨流體運動（Alfvén, 1942）。由于??太陽風正是這種流體??，??太陽的磁力線被凍結在太陽風中??，并??被迫隨太陽風傳播??（Roberts, 2007），從而形成??行星際磁場（IMF）??。

根據??起源位置??的不同——無論是來自??日冕洞??還是??太陽赤道帶??，太陽風可分為??速度達750 km/s的快太陽風??和約400 km/s的慢太陽風??（Feldman等, 2005）。??日冕洞在X射線圖像中呈現暗區??，因為該區域??溫度更低、密度更稀薄??（Cranmer, 2009），且??磁場不會回折而是延伸至行星際空間??（Parker, 1959），因此??等離子體可自由流出形成快太陽風??。??在太陽活動極大期，日冕洞可能出現在日冕任何區域??；而??在太陽活動極小期，它們通常出現在太陽兩極?（McComas等, 2003）。

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Figure 3.3: Combined X-ray image of the Sun’s active regions observed from several telescopes. High-energy X-rays from NASA’s Nuclear Spectroscopic Telescope Array (NuSTAR) are shown in blue, low-energy X-rays from Japan’s Hinode spacecraft are green, and extreme ultraviolet light from NASA’s Solar Dynamics Observatory (SDO) is yellow and red (Credit: NASA/JPL-Caltech/GSFC/JAXA, 2015)

The expansion of the solar wind is a radial propagation away from the Sun?? that it is ??diluted and cools down on its journey??. As seen from some observations the ??density of the solar wind decreases approximately as r^(-2)?? while the ??temperature decrease is less significant about a factor of 20?? (Baumjohann and Treumann, 2012). During the expansion of the radial flow, the ??Sun rotates, which is 27 days on average??, and because the ??magnetic field lines are anchored in the sun??, the ??radial flow looks like an Archimedean spiral called the Parker Spiral??. At ??1 AU the spiral makes 45° to the Earth-Sun line??, see ??Figure 3.4??.

??太陽風的膨脹是遠離太陽的徑向傳播??，在此過程中??逐漸稀釋冷卻??。觀測表明，??太陽風密度近似按r^(-2)遞減??，而??溫度下降較緩，降幅約20倍??（Baumjohann和Treumann, 2012）。由于??太陽自轉（平均周期27天）??且??磁力線錨定于太陽??，??徑向流會形成阿基米德螺旋線結構（帕克螺旋）??。在??1天文單位（AU）處??，該螺旋與日地連線呈??45°夾角??，參見??圖3.4??。

Figure 3.4: Schematic representation of the Archimedes spiral structure of the interplan etary magnetic field (Figure is from Bittencourt, 2004; Figure 3)

The solar wind travels throughout the solar system and defines the heliosphere?? the region in space whose frontier is impeded by the interstellar medium. Consequently, ??the heliosphere is a giant bubble whose center is the Sun?? and protects the solar system from ??interstellar radiations and cosmic rays??. The size of the heliosphere is about ??121 AU?? (Cowen, 2013). The ??solar wind decelerates when it flows outward through the Solar System??. At a certain point, ??the flow becomes subsonic and forms a shock??. Its name is ??termination shock?? (Jokipii, 2013). Then the flow continues its journey outward and ??interacts with the interstellar material??. It is under debate whether ??the speed of the Solar System is super- or sub-sonic??. A ??bow shock or bow wave forms before the heliosphere??, respectively. Furthermore, due to the interaction, ??the solar wind becomes turbulent?? and this region is called ??heliosheath?? which is in between the ??termination shock and the heliopause?? (Burlaga, 2015).

太陽風在太陽系中傳播并定義了日球層??——這個??空間區域的邊界受星際介質阻擋??。因此，??日球層是一個以太陽為中心的巨型氣泡??，它??保護太陽系免受星際輻射和宇宙射線的侵害??。日球層的尺寸約為??121天文單位（AU）??（Cowen，2013）。??當太陽風向外穿過太陽系時會逐漸減速??。在某個臨界點，??太陽風流會變為亞音速并形成激波這個激波被稱為??終止激波??（Jokipii，2013）。隨后，太陽風繼續向外傳播并??與星際物質相互作用??。目前學界仍在爭論系的速度是超音速還是亞音速??的問題，這分別會導致在日球層前形成??弓激波或弓形波??。此外，由于這種相互作用太陽風會變得湍流??，這個區域被稱為??日球鞘??，它位于??終止激波和日球層頂之間?**?（Burlaga，2015）。

?▍3.3 Interplanetary shocks?▍

The average interplanetary magnetic field strength is around 6 nT at 1 AU?? (Lowrie and Fichtner, 2020) and considering the average speed of the solar wind, ??the solar wind is both supersonic and super-Alfvénic??, which means it is ??super-magnetosonic??. As a result of this nature of the solar wind, when it ??collides with celestial objects?? such as planets, moons, asteroids, and comets or its ??slow-flowing part, shocks are formed??. The main characteristic of shocks?? that can be found in interplanetary space to interstellar regions is ??a denser state in contrast to the medium?? in which they propagate due to the shock formation. Furthermore, if the ??celestial objects are magnetized??, then the ??shock formations create interesting interaction regions?, and their physics looks very exciting.

??3.3 行星際激波?

??在1天文單位（AU）處，平均行星際磁場強度約為6納特斯拉（nT）??（Lowrie和Fichtner，2020）。考慮到太陽風的平均速度，??太陽風既是超音速的，又是超阿爾芬速的??，這意味著它是??超磁聲速的??。由于太陽風的這種特性，當它??與行星、衛星、小行星和彗星等天體或其低速部分碰撞時??，就會??形成激波??。在行星際空間到星際區域中發現的激波，其??主要特征是由于激波形成，它們傳播的介質會呈現更致密的狀態??。此外，如果??天體被磁化??，那么??激波的形成會產生有趣的相互作用區域??。

One such fascinating magnetized object is our planet, the Earth.?? The Earth has its ??magnetic field, known as the magnetosphere??, which originated from its ??inner core via the dynamo effect?? (Gilbert, 2003). ??The Earth's magnetosphere extends 60000 kilometers in the sunward direction while a million kilometers in the anti-sunward direction?? (Lakhina et al., 2009). Such a ??big extension is a common characteristic of all the magnetized planets??, and as a result of this extension, ??the cross-section of a planet is increased by a large factor??. For example for the Earth, ??the factor is 150?? (Baumjohann and Treumann, 2012). ??

The magnetic field frozen in the supersonic and super-Alfvénic solar wind plasma cannot enter the magnetosphere??, the region where the ??terrestrial magnetosphere dominates??. As a result, a ??special standing wave, the bow shock forms (Baumjohann and Nakamura, 2007). ??Consequence of the slowing down of the solar wind plasma??, the ??kinetic energy of some of the particles is converted into thermal energy??, which occurs behind the bow shock, and this region is called the ??magnetosheath??. The ??boundary region between two magnetic field lines is called the magnetopause??. ??Bow shock formation is common to all the planets and celestial objects with or without magnetospheres?? (Mazelle et al., 2004).

地球??就是??一個磁化天體??。地球擁有自身的??磁場，稱為磁層??，其源于地球??內核的發電機效應??（Gilbert，2003）。??地球磁層向太陽方向延伸60000公里，而背向太陽方向延伸達100萬公里??（Lakhina等，2009）。這種??巨大延伸是所有磁化行星的共同特征??，由于這種延伸，??行星的橫截面積被顯著放大??。以地球為例，??放大因子達150倍（B??umjohann和Treumann，2012）。

??凍結在超音速和超阿爾芬速太陽風等離子體中的磁場無法進入磁層??——這個??地球磁場主導的區域??。因此會形成一種??特殊的駐波，即弓激波??（Baumjohann和Nakamura，2007）。由于??太陽風等離子體減速??，??部分粒子的動能轉化為熱能??，這一過程發生在弓激波后方，該區域稱為??磁鞘??。??兩條磁力線之間的邊界區域稱為磁層頂??。??弓激波的形成對所有具有或不具有磁層的行星和天體都是普遍現象??（Mazelle等，2004）。

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Figure 3.5: Illustration of the Earth’s magnetosphere and its interaction with the solar wind. (Figure is from Kivelson and Bagenal, 2007; Figure 1)

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Figure 3.6: Illustration of a CME event. (Credit: NASA, 2017)

In addition, there are several other shocks??, the previously mentioned ??termination shock??, ??coronal mass ejection (CME) driven shock??, and a ??co-rotating interaction region (CIR) driven shock??. When a ??CME event occurs??, it ??moves faster than the background solar wind flow??, resulting in a ??shock wave??. On the created shock waves ??charged particles accelerate??. So usually ??CMEs are one of the main causes of the solar energetic particles (SEPs)?? (Cane and Lario, 2006), see ??Figure 3.6??.

??Similarly to the shocks caused by CMEs??, When the ??fast-moving solar wind flow catches the slow-moving solar wind flow??, the so-called ??co-rotating interaction region is formed??, and if the ??pressure gradient gets sufficiently large?? and the ??speed difference surpasses the local speed shocks can arise?? (Heber et al., 1999), ??Figure 3.7??.

??此外還存在其他幾種激波??，前文提到的??終止激波??、??日冕物質拋射（CME）驅動的激波??，以及??共轉相互作用區（CIR）波??。當??CME事件發生時??，其運動速度??超過背景太陽風流速??，從而產生??激波??。在這些激波上，??帶電粒子被加速??。因此??CME通常是太陽高能粒子（SEPs）的主要來源之一??（Cane和Lario，2006），參見??圖3.6??。

??與CME產生的激波類似??，當??快速運動的太陽風流追上慢速太陽風流??時，會形成所謂的??共轉相互作用區??；若??壓力梯度足夠大且速度差超過當地聲速??，就可能產生激波（Heber等，1999），??圖3.7??。

Figure 3.7: Co-rotating interaction region. (Figure is from David Burgess, 2017)

3.3.1 Classifications of Interplanetary (IP) shocks

Figure 3.8: Categorizations of IP shocks. N, T, B, and V denote number density, the proton temperature, the magnitude of the magnetic field, and speed respectively. (Figure is from WIND MFI Team Science, 2001)

Armed with the above shock definitions and categorizations knowledge??, now we can finally be able to ??classify IP shocks??. The ??IP shocks can be classified based on their travel directions concerning the solar wind frame of reference??. If the shock is ??moving away from its driver??, here the Sun is referred but ??drivers can be detailed such as ICMEs, CIRs, etc??, the shock is called ??forward shock (FS)??, and if it is ??moving toward its driver??, it is called ??reverse shock (RS)??. Adding the previous definitions of ??fast and slow shocks??, IP shocks are usually categorized as ??fast forward (FF), fast reverse (FR), slow forward (SF) and slow reverse (SR) shocks?? (Berdichevsky et al., 2000), see ??Figure 3.8??.

基于上述激波定義和分類知識??，我們現在終于能夠??對行星際激波（IP shocks）進行分類??。行星際激波可根據其??相對于太陽風參考系的傳播方向??進行分類。若激波??遠離其驅動源（此處指太陽，但驅動源可具體分為ICMEs、CIRs等）??，則稱為??前向激波（FS）??；若激波??朝向其驅動源傳播??，則稱為??反向激波（RS）結合先前對快慢激波的定義，行星際激波通常被分類為??快前向（FF）、快反向（FR）、慢前向（SF）和慢反向（SR）激波??（Berdichevsky等，2000），參見??圖3.8?。

As you can see from ??Figure 3.8??, the ??solar wind plasma parameters?? – ??number density N, the proton temperature T, the magnitude of the magnetic field B, and the bulk speed V parameters increase dramatically from upstream (unshocked) to downstream (shocked) regions in fast forward (FF) shocks?? while the ??parameters except for the bulk speed decrease in fast reverse (FR) shocks?? [1]. In the case of ??slow forward (SF) shocks??, the ??parameters except for the magnitude of the magnetic field increase from upstream to downstream?? whereas in the case of ??slow reverse (SR) shocks??, the ??number density N and the proton temperature T decrease while the magnitude of magnetic field B and the bulk speed V increase?? [2].

??Within 1 AU, the most frequent IP shocks are fast forward (FF) shocks?? (Richter et al., 1985).

[1]???In a sense of the reverse upstream to downstream in fast reverse (FR) shocks??, the ??parameters except for the bulk speed increase from reversed upstream to downstream??

[2]???Again in the sense of the reversed upstream to downstream in slow reverse (SR) shocks??, the ??number density N and the proton temperature T increase while the magnitude of magnetic field B and the bulk speed V decrease?

如??圖3.8??所示，在??快前向（FF）激波??中，??太陽風等離子體參數??——??數密度N、質子溫度T、磁場強度B和整體速度V??從上游（未受擾）到下游（受擾）區域??顯著增加??；而在??快反向（FR）激波??中，??除整體速度外其他參數均下降 ??[1]。對于??慢前向（SF）激波??，??除磁場強度外其他參數均隨上游至下游方向增加??；而對于??慢反向（SR）激波??，??數密度N和質子溫度T下降，而磁場強度B和整體速度V上升????[2]。