Space Physics and Aeronomy, Volume 1, Solar Physics and Solar Wind E-Book

• Herausgeber: John Wiley & Sons
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Beschreibung

A comprehensive view of our Sun at the start of a new era in solar and heliospheric physics Humans have been observing and studying our Sun for centuries, yet much is still unknown about the processes that drive its behavior. Thanks to a new generation of space missions and ground telescopes, we are poised to dramatically increase our understanding of the Sun and its environment. Solar Physics and Solar Wind explores advances in solar and heliospheric research over recent decades, as well as the challenges that remain. This comprehensive reference work covers the solar interior, magnetism and radiation, plasma heating and acceleration, the sun's atmosphere, and solar activity. Volume highlights include: * Explanations for processes in the solar interior * New insights on the solar wind * The challenges of measuring the Sun's magnetic field and its radiative output * Description of solar atmospheric phenomena such as spicules and jets * New developments in understanding flares and coronal mass ejections * Ongoing research into how the solar corona is heated The American Geophysical Union promotes discovery in Earth and space science for the benefit of humanity. Its publications disseminate scientific knowledge and provide resources for researchers, students, and professionals. Find out more about the Space Physics and Aeronomy collection in this Q&A with the Editors in Chief

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258  Solar Physics and Solar WindNour E. Raouafi and Angelos Vourlidas (Eds.)

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Space Physics and Aeronomy Collection Volume 1Geophysical Monograph 258

Solar Physics and Solar Wind

Nour E. Raouafi

Angelos VourlidasEditors

Yongliang Zhang

Larry J. PaxtonCollection Editors in Chief

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Library of Congress Cataloging‐in‐Publication Data

Names: Raouafi, Nour E., editor. | Vourlidas, Angelos, editor.Title: Solar physics and solar wind / Nour E. Raouafi, Angelos Vourlidas, editors.Description: Hoboken, NJ : Wiley, 2021. | Series: Geophysical monograph series | Includes bibliographical references and index.Identifiers: LCCN 2020047581 | ISBN 9781119507536 (hardback) | ISBN 9781119815488 (adobe pdf) | ISBN9781119815471 (epub)Subjects: LCSH: Sun. | Solar wind.Classification: LCC QB521 .A88 2021 | DDC 523.7–dc23LC record available at https://lccn.loc.gov/2020047581

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LIST OF CONTRIBUTORS

Olga AlexandrovaLaboratoire d’Etudes Spatiales et d’Instrumentation en AstrophysiqueObservatoire de ParisUniversité PSLCNRSSorbonne UniversitéUniversité de ParisMeudon, France

Alessandro BemporadINAF Osservatorio Astrofisico di TorinoTurin, Italy

Luca BertelloNational Solar ObservatoryBoulder, Colorado, USA

Christina M. S. CohenSpace Radiation LaboratoryCalifornia Institute of TechnologyPasadena, California, USA

Serena CriscuoliNational Solar ObservatoryBoulder, Colorado, USA

Mausumi DikpatiHigh Altitude ObservatoryBoulder, Colorado, USA

Cooper DownsPredictive Science Inc.San Diego, California, USA

Aleida K. HigginsonNASA Goddard Space Flight CenterGreenbelt, Maryland, USA

Rachel HoweSchool of Physics and AstronomyUniversity of BirminghamBirmingham, United Kingdom

James A. KlimchukNASA Goddard Space Flight CenterGreenbelt, Maryland, USA

Michael LavarraInstitut de Recherche en Astrophysique et PlanétologieToulouse, France

Benoit LavraudLaboratoire d’Astrophysique de BordeauxUniversité de BordeauxCNRS, B18NPessac, France

Gang LiDepartment of Space ScienceUniversity of Alabama in HuntsvilleHuntsville, Alabama, USA

Mark LintonSpace Science DivisionNaval Research LaboratoryWashington, D.C., USA

Noé LugazUniversity of New HampshireDurham, New Hampshire, USA

Glenn M. MasonJohns Hopkins University Applied Physics LaboratoryLaurel, Maryland, USA

Lorenzo MatteiniDepartment of PhysicsImperial College LondonLondon, UK

Ineke De MoortelSchool of Mathematics and StatisticsUniversity of St AndrewsSt Andrews, United Kingdom;andRosseland Centre for Solar Physics University of OsloNorway

Susanna ParentiUniversité Paris‐SaclayCNRSInstitut d’Astrophysique SpatialeOrsay, France

Gordon PetrieNational Solar ObservatoryBoulder, Colorado, USA

Viviane PierrardBelgian Institute for Space AeronomyBrussels, Belgium

Rui PintoLaboratoire Dynamique des Etoiles, des (Exo)planètes et de leur Environnement (LDE3)Astrophysics Division (DAp/AIM)Saclay Nuclear Research Centre (CEA Saclay)Gif‐sur‐Yvette, France;andInstitut de Recherche en Astrophysique et PlanétologieToulouse, France

Jiong QiuMontana State UniversityBozeman, Montana, USA

Nour E. RaouafiJohns Hopkins University Applied Physics LaboratoryLaurel, Maryland, USA

Fabio RealeDipartimento di Fisica & ChimicaUniversitá di PalermoPalermo, Italy

Alexis P. RouillardInstitut de Recherche en Astrophysique et PlanétologieToulouse, France

Eduardo Sanchez‐DiazInstitut de Recherche en Astrophysique et PlanétologieToulouse, France

Albert Y. ShihNASA Goddard Space Flight CenterGreenbelt, Maryland, USA

Alphonse C. SterlingNASA Marshall Space Flight CenterHuntsville, Alabama, USA

Barbara J. ThompsonNASA Goddard Space Flight CenterGreenbelt, Maryland, USA

Nicholeen ViallNASA Goddard Space Flight CenterGreenbelt, Maryland, USA

Christian VocksLeibniz Institute for Astrophysics PotsdamPotsdam, Germany

Angelos VourlidasJohns Hopkins University Applied Physics LaboratoryLaurel, Maryland, USA

Linghua WangSchool of Earth and Space SciencesPeking UniversityBeijing, China

David F. WebbInstitute for Scientific Research Boston College Chestnut Hill, Massachusetts, USA

Yihong WuLeibniz Institute for Astrophysics PotsdamPotsdam, Germany

PREFACE

The upcoming decade will mark a turning point in solar and heliospheric research. The spacecraft fleet comprising the Heliophysics Systems Observatory will be augmented by a set of unique space missions and ground‐based telescopes that will greatly expand our knowledge of the heliosphere. The Parker Solar Probe (PSP) was launched in August 2018, sixty years after it was conceived as humanity’s first mission to enter and study a stellar atmosphere from “within.” The PSP will approach closer to the Sun than any spacecraft before to explore the solar atmosphere as close as 8.86 solar radii above the surface. The PSP has completed 6 of its 24 scheduled orbits as of this writing. The data recorded during the first six orbits show an unprecedented view of the nascent solar wind. We expect more discoveries as the spacecraft flies ever closer to the Sun. The Solar Orbiter (SolO), launched in February 2020, will approach within 60 solar radii but from an orbit inclined by up to 34° out of the ecliptic. SolO will give us the first unprojected images of the solar poles and measure the magnetic fields in these regions. The largest ground‐based telescope ever built, the 4‐m Daniel K. Inouye Solar Telescope (DKIST), will be operational next year. The DKIST will reveal solar structures as small as 20–30 km in diameter. In the near future, the 4‐m European Solar Telescope (EST) will also become operational.

The observations and measurements from the PSP and SolO and those we will get from the DKIST, EST, and other future missions and telescopes will not only help answer long‐standing scientific questions but also lead to significant discoveries and open new avenues for exploration. These measurement capabilities will write a new chapter of space and solar physics. So, it is a good time now to take a look at the status of the solar and heliospheric research. This book presents seven chapters that cover most aspects of solar and heliospheric physics.

Important technological inventions, along with essential advances in mathematics and physical theories made over the last few centuries, led to fundamental solar and astronomical discoveries. These discoveries created new puzzles that form the main pillars and axes of solar and heliospheric research. The invention of the telescope and the discovery of sunspots by Galileo in the 16th century represent a turning point in solar physics research. Soon after, the 11‐year solar cycle was revealed with its long‐ and short‐term variability. The observation of the strongest flare in recorded history by Carrington in 1858 was another important milestone and the building block for solar activity and space weather. The exact nature of these phenomena remained hidden until the discovery of solar magnetism by George Hale (1908). He stated in his The Astrophysical Journal paper: “The present paper describes an attempt to enter one of the new fields of research opened by this recent work with the spectroheliograph,” which is a testament to the fact that scientific advances go hand in hand with technology. In the next half‐ century, three major phenomena were discovered: the coronal heating problem, the solar wind and its acceleration, and coronal mass ejections. All of these phenomena are driven by the magnetic field, which is generated deep in the convection zone

The Sun remains the most observed and studied star in the Universe. Yet many of the processes that govern its behavior are not fully understood. The solar dynamo and convection, coronal heating, the acceleration of the solar wind, flares, coronal mass ejections, and solar energetic particles are all outstanding examples of challenging solar and heliospheric problems. Solar and heliospheric research has undergone a renaissance in the last decades. Since the advent of the Space Age, significant strides have been made in our understanding of the Sun along with breakthrough discoveries. Important advances in theory and computing power are keeping pace, allowing access to fundamental physical processes that are otherwise impossible to explore because of the complexity of the underlying theories. These advances have helped us better understand the challenges facing us when we try to comprehend how the Sun and its corona work and also what it takes to obtain the measurements needed to make progress. Figure I.1 is an illustration of the integrated ground–space–theory system, which is a prerequisite to overcoming the hurdles we were facing for decades.

Understanding the Sun is essential not only because we live in its extended atmosphere (i.e., the corona and the solar wind) but also because it is the only star we can study in detail. The knowledge we gain from observing the Sun and its environment provides insights into other worlds that may harbor life like our Earth.

This book provides an overview of solar physics and the advances made over the past few decades as well as the challenging problems that remain. The seven chapters cover the solar interior, the atmosphere, magnetism and radiation, plasma heating and acceleration, and solar activity. It is a comprehensive view of how our star works and what is needed to understand it better. We hope that this reference work will help researchers in other fields, young scientists, teachers, students, and the public to familiarize themselves with the status of the field as of 2019.

Nour E. Raouafi and Angelos VourlidasJohns Hopkins University Applied Physics LaboratoryUSA

Figure I.1 Illustration of the complex solar environment and the recent observational and modeling advances that will lead to breakthrough insights into challenging phenomena that have been puzzling scientists for decades. Top left: Slice view showing the Sun at different wavelengths. Top right: The fleet of heliospheric space missions. Bottom left: Image of the 4‐m DKIST solar telescope. Bottom right: Simulation of the structure of the solar corona for the 2017 eclipse.

(Source: Reproduced with permission from Predictive Science Inc.)

1The Solar Wind

Alexis P. Rouillard1, Nicholeen Viall2, Viviane Pierrard3, Christian Vocks4, Lorenzo Matteini5, Olga Alexandrova6, Aleida K. Higginson2, Benoit Lavraud7, Michael Lavarra1, Yihong Wu4, Rui Pinto1,8, Alessandro Bemporad9, and Eduardo Sanchez‐Diaz1

1 Institut de Recherche en Astrophysique et Planétologie, Toulouse, France

2 NASA Goddard Space Flight Center, Greenbelt, Maryland, USA

3 Belgian Institute for Space Aeronomy, Brussels, Belgium

4 Leibniz Institute for Astrophysics Potsdam, Potsdam, Germany

5 Department of Physics, Imperial College London, London, UK

6 Laboratoire d’Etudes Spatiales et d’Instrumentation en Astrophysique, Observatoire de Paris, Université PSL, CNRS Sorbonne Université, Université de Paris, Meudon, France

7 Laboratoire d’Astrophysique de Bordeaux, Université de Bordeaux, CNRS, B18N, Pessac, France

8 Laboratoire Dynamique des Etoiles, des (Exo)planètes et de leur Environnement (LDE3), Astrophysics Division (DAp/ AIM), Saclay Nuclear Research Centre (CEA Saclay), Gif‐sur‐Yvette, France

9 INAF Osservatorio Astrofisico di Torino, Turin, Italy

1.1. INTRODUCTION

One of the current mysteries in heliophysics is the heating of the solar atmosphere to temperatures that are orders of magnitude hotter than the solar surface. As a result of this heating, the Sun cannot contain its atmosphere, and a continual outflow of plasma streams out from the solar corona to interplanetary space and beyond. For the debate surrounding the exact physical mechanisms of the heating of the corona, we direct the reader to Chapter 2. We here discuss the physical mechanisms behind the formation and propagation of the solar wind that are not yet well understood.

The whole volume of space influenced by the solar wind is called the heliosphere, and its size is in part modulated by the solar wind ram pressure. The solar wind extends from the corona to well beyond a hundred astronomical units (AU) to a termination shock. Beyond that shock, the solar wind slows down abruptly in response to the pressure of the interstellar medium, and the plasma becomes compressed and more turbulent until it reaches a zone where it can no longer push back the interstellar plasma. In situ measurements of these outer regions are sparse, and our exploration of this boundary layer has only just begun with the Voyager spacecraft. In particular, the global shape of the heliosphere is still not known because the Voyager has only measured a very small region of the heliospheric boundary.

Birkeland (1908) argued very early that a corpuscular emission from sunspots consisting of relativistic electrons must impact Earth’s magnetic field and be deflected to the polar regions to create the aurora. For several decades, it was realized that particles could be emitted from the Sun during flares, but it was generally thought that the space around Earth was mostly empty or perhaps traversed by occasional streams of particles from the Sun (Chapman & Ferraro, 1931). The prevailing view was that the solar corona consists of a hot gas (possibly extending to 1 AU), in thermal and hydrostatic equilibrium, pulled back by the solar gravitational field (Chapman & Zirin, 1957). Detailed observational studies of comets by Biermann (1951) showed that a subset of their tails cannot be accelerated by radiation pressure alone but may also respond to material flowing out from the Sun’s atmosphere. He suggested that the passage of solar particles at the comet formed an ion tail and that these particles must have a very high speed relative to the comet in order to align the tail in the Sun’s direction. Parker (1958) built on these observations and realized that the high temperature of the corona can provide enough energy to force coronal plasma to accelerate from subsonic to supersonic speeds. He demonstrated that the hydrostatic approach predicted too high kinetic pressure at infinity and that a continuous radial expansion of solar gas must act to reduce the coronal pressure. This was the first theory describing the continual expansion of what we now call the solar wind.

In this model, a dominant force affecting coronal particles and pushing them outward is induced by the thermal pressure gradient in the corona. Parker’s original model assumed an isothermal corona, but subsequent models allowing for a varying temperature with radial distance confirmed that a supersonic wind can also form under such conditions. The presence of an outflowing supersonic wind ranging from 300 to 800 km/s was confirmed by plasma measurements from the Luna 1, 2, 3; Venera 1; and Mariner 2 (e.g., Neugebauer & Snyder, 1962) and the numerous subsequent solar wind dedicated missions (e.g., Hundhausen & Gosling, 1976; Marsch et al., 1982). The Parker model provided an acceleration from subsonic speed in the coronal region to supersonic speeds of around 300–400 km/s typical of the slow solar wind. His model was unable, however, to explain fast solar wind speeds of 700–800 km/s without assuming unrealistic coronal hole temperatures in excess of 2 × 106K. As we shall show in this chapter, coronal and solar wind measurements suggest that additional physical mechanisms must be accounted for to produce a fast solar wind.

The solar wind has been measured in situ for several decades near 1 AU. Typical solar wind speeds near 1 AU range from 300 to 800 km/s, proton temperatures take values between 105 K in the slow wind to about 2 × 105 in the fast wind. There is now no doubt that the fast solar wind measured in situ originates near the center of coronal holes at the Sun (see Chapters 2 and 3). Coronal holes are cooler regions of the solar atmosphere that exhibit drops in Extreme UltraViolet (EUV) emissions. The cooler temperatures result from the significant escape of heat out in the solar wind. The fast solar wind streams out along open magnetic fields connecting regions deep inside coronal holes to the interplanetary medium. The origin of the slow wind is more complex and likely consists of multiple source components: these include transient releases from helmet streamers, plasma accelerating on rapidly expanding magnetic fields rooted at the boundary of coronal holes (through processes likely similar to that of the fast wind), and/or continual plasma exchanges between loops and open magnetic fields. We will describe these different components in the following sections.

1.2. OBSERVATIONS OF THE NASCENT SOLAR WIND

1.2.1. Remote‐Sensing Observations of Coronal Heating and the Solar Wind

We begin our story on the solar wind with its formation in the solar corona; we first present remote‐sensing observations that have provided important information on the conditions in which the winds are produced.

It is hard to observe the coronal source regions of the fast and slow solar winds in white‐light images obtained routinely by coronagraphs. This is because the fast solar wind originates in very tenuous coronal holes, whereas the slow solar wind emerges from the vicinity of dense streamer loops that completely dominate coronal brightness. In contrast, spectroscopic observations provide more detailed information about the temperatures, flow velocities, and wave properties during the formation of the winds near their sources.

The Solar and Heliospheric Observatory (SoHO; Domingo et al., 1995) revolutionized how we observe the corona and the nascent solar wind. In particular, the Ultraviolet Coronagraph Spectrometer (UVCS/SoHO; Kohl et al., 1995) revealed that heavy ions such as O5+ and Mg9+ are heated hundreds of times more strongly than protons and electrons, and have very anisotropic kinetic temperatures (Antonucci et al., 2000; Kohl et al., 1997, 1998), meaning that temperatures measured in the direction perpendicular to the magnetic field are often much larger than those parallel to the field (see Chapter 2). A comparison between different ion and electron temperatures derived from coronal observations and solar wind measurements are shown in Figure 1.1. The measured temperatures provide support to coronal heating mechanisms involving collisionless wave–particle resonances with frequencies of 10 Hz to 10 kHz. These waves could be damped to heat preferentially heavy ions (Cranmer et al., 1999; Tu & Marsch, 1997). It is not clear yet how such hypothetical waves can be generated from lower‐frequency Alfvén waves typically emitted at minute periods. Some proposed mechanisms involve the turbulent cascade of magneto‐plasma fluctuations from low to high frequencies (Hollweg, 2002). Then kinetic processes must occur to damp the small‐scale fluctuations that have spectrally cascaded to oblique wavevectors. Proposed mechanisms include ion‐cyclotron and Landau damping (Leamon et al., 1998). Nonlinear processes such as beam instabilities or mode conversion and damping have also been proposed.

Figure 1.1 Radial evolution of solar wind temperatures from the corona to 1 AU. Indirect estimates of coronal temperatures are derived for three species (electrons, oxygen, and hydrogen) from an “empirical coronal model” that exploits SUMER (1–1.2R⊙) and UVCS (1.5–4R⊙) line widths compared with direct in situ measurements at distances greater than 60R⊙) in the high‐speed wind (Cranmer et al., 1999). The in situ data were assembled from Helios, IMP, Ulysses, and Voyager particle data, and double sets of curves denote rough lower and upper bounds on representative fast‐wind values. The upper and lower limits correspond to extreme values of an assumed non‐thermal component to the line broadening that is attributed to unresolved MHD wave motions along the line of sight. Temperatures of electrons (solid black), hydrogen (dotted), and oxygen (dashed) are shown. Oxygen ions correspond to O5+ in the corona but O6+ in the far solar wind, and coronal temperatures of neutral hydrogen are here compared with proton temperatures in the solar wind. Figure taken with permission from (Cranmer et al., 1999).

Emission lines, and more specifically the dimming of certain lines measured via coronal spectroscopy, have been used to infer the outflow speed and density of the forming fast solar wind. This effect is most important for spectral lines in the UV range that have a significant component due to ion excitation via resonant scattering of chromospheric emission (followed by spontaneous emission). For some lines, such as O VI doublet 103.2 nm and 103.8 nm lines, and of the HI Lyman‐ α 121.6 nm and Lyman‐ β 102.5 nm lines, the collisional and radiative components can be separated (Marocchi et al., 2001). This separation can be used to investigate the dimming of the radiative component, which is due to the radial expansion of the emitting coronal atoms. This Doppler shifts the (narrow) exciting chromospheric profile with respect to the (broad) atomic absorption profile. As a result, the UV intensity of the resonantly scattered component of the line emission decreases with increasing outflow velocities (Hyder & Lites, 1970; Noci et al., 1987; Withbroe et al., 1982). These observations applied to the O VI doublet lines have shown that the fast solar wind becomes supersonic much closer to the Sun than the slow solar wind. The fast O5+ ions reach speeds in excess of 600 km/s within 4 solar radii from the solar surface (Antonucci et al., 2000). Identical techniques were also used to study flows in the vicinity of streamers, adjacent but not above the helmet, and found that the slow solar wind accelerates more slowly, with its outflow speed remaining below 200 km/s at least until 4 solar radii (Abbo et al., 2010; Strachan et al., 2000). Figure 1.2 presents a summary of these results.

Figure 1.2 Outflow velocity (km/s) of the solar wind for the considered four regions as a function of the heliocentric distance (in solar radii). The gray band from 4 to 10 R⊙ shows the range of outflow velocities for the slow wind obtained with LASCO (Sheeley et al., 1997). The solid curve up to 3 R⊙ represents the values of the fast wind obtained from the UVCS data (Antonucci et al., 2000), and the dashed curves show the results by Telloni et al. (2007) of the fast wind velocity. The error bars (small in many cases) have been estimated based on the propagation of the statistical uncertainties of the observed OVI 1032 and 1037 line intensities.

(Source: Reproduced from Abbo et al., 2010. © 2010, Elsevier.)

A similar analysis of the Doppler dimming technique has been applied to the Lyman‐ α (121.6 nm) emission of neutral hydrogen. The latter moves outward in response to rapid charge‐exchange coupling with the heated protons that eventually form the bulk of the solar wind. Neutral hydrogen therefore acts as a proxy for protons at heights up to 2.5 solar radii in coronal holes, and higher heights in more dense structures. In some coronal structures, He+ can act as a proxy for alpha particles, which are also an important component of the solar corona. The UVCS/SoHO instrument has provided H I Lyman‐ α spectral line data over 18 years. This makes it possible to study proton outflow speeds throughout the solar cycle, focusing on the coronal region sampled by the spectrometer field of view, such as coronal streamers (Susino et al., 2008; Zangrilli & Poletto, 2016) and coronal holes (Antonucci et al., 2000; Strachan et al., 1993; Teriaca et al., 2003).

In a recent analysis (Bemporad, 2017), UVCS daily Lyman‐ α synoptic data were combined to provide the first 2D images of coronal Lyman‐ α emission, representative of future data that will be acquired by the Metis coronagraph onboard Solar Orbiter (Antonucci et al., 2017). These have been directly combined with classical 2D coronagraphic images acquired in white light with LASCO to derive 2D maps of HI outflow speeds, with a technique originally described by Withbroe et al. (1982) that neglects line‐of‐sight integration effects. As pointed out by Bemporad (2017), because both the radiative component of Lyman‐ α emission and the white‐light polarized emission depend on the electron density distribution integrated along the line of sight, this latter quantity can be simplified by directly taking the ratio between the two UV and white‐light intensities.

Figure 1.3 shows an example of a 2D outflow velocity map, ranging between 1.5 and 4.0 R⊙, that was obtained this way (Bemporad, 2017). The map shows an increasing speed with altitude from about 150–200 km/s in the equatorial regions to 400 km/s in the polar regions. As we shall see, these values are in agreement with the expected latitudinal distribution of slow and fast solar wind components during solar minimum, corresponding to equatorial regions, and mid‐latitude and polar regions, respectively. Alternatively, as shown more recently by Dolei et al. (2018), line‐of‐sight integration effects can be fully taken into account (under some assumptions) by deriving electron densities from white‐light coronagraphic images. From this, one can derive HI outflow speed maps by again exploiting the Doppler dimming technique discussed above.

Figure 1.3 2D map of radial outflow velocity in the plane of the sky derived from the ratio between white‐light and UV coronal emissions. The outer white region corresponds to altitudes where the Doppler dimming technique with the Ly spectral line cannot be applied anymore.

(Source: Taken from Bemporad, 2017. © 2017, IOP Publishing.)

The solar wind acceleration has also been measured via oscillations of the coronal plasma, identified as Alfvén waves in coronagraphic observations. Alfvén waves transverse to the plane‐of‐sky have been observed to be omnipresent in coronal holes within the first 0.3R⊙ of the solar atmosphere by the Coronal Multichannel Polarimeter (Tomczyk et al., 2016). Detailed analysis of off‐limb Doppler shifts (Fe XIII) at high time cadence has furthermore revealed the presence of upward‐ and downward‐propagating waves and, by comparing their respective phase speeds, allowed for the determination of the bulk flow speed profile of nascent fast wind flows (Morton et al., 2015, 2016). These observations have furthermore shown that the spectra of the oscillations show a predominant 1/f slope, reminiscent of solar wind measurements made beyond 0.3 AU by past space probes (e.g., Bavassano, Dobrowolny, Fanfoni et al., 1982; M. L. Goldstein et al., 1995). This supports the idea that the 1/f component of the oscillation spectra observed in the solar wind are already set in the low solar atmosphere and are advected outward (see, for example, Verdini et al., 2012). The statistical properties of fluctuations at magnetohydrodynamic (MHD) and kinetic scales are discussed further in Section 1.4.

1.2.2. Transient Coronal Outflows in the Nascent Solar Wind

The formation of the background solar wind, introduced in the previous section, is continually perturbed by the ejection of jets and small transients that form in the corona. Direct observations of these transient outflows in EUV and white‐light images by the STEREO and SoHO spacecraft have provided new insights on the origin of mesoscale structures measured in situ in the solar wind.

Variable solar wind outflows in the form of plasmoids are continually released from helmet streamers in white‐light (i.e., electron density) observations (e.g., Harrison et al., 2009; Rouillard, Davies, et al., 2010; Rouillard, Lavraud, et al., 2010; Rouillard et al., 2009; Sheeley et al., 1997; Sheeley et al., 2007; Wang et al., 1998; Wang et al., 2000). These plasmoids have been tracked from the corona, through the inner heliosphere, and in some cases out to 1 AU using heliospheric imagers (Rouillard, Davies, et al., 2010; Rouillard et al., 2011; Sheeley & Rouillard, 2010), showing in a direct way that some of the helmet streamer structures produced in the corona result in density structures measured in the inner heliosphere. Plasmoids (or “blobs”) have been tracked from the tip of streamers where they typically form to several tens of solar radii, and analysis of their kinematic properties has confirmed that they are advected in the slow wind (Sheeley et al., 1997). In fact, this type of analysis has provided one of the rare kinematic measurements of the forming slow wind. It has revealed that a subset of the slow solar wind is released right above helmet streamers and accelerates over 20–30 solar radii to reach its terminal speed of about 300 km/s (this acceleration is shown as the gray area in Figure 1.2).

Helmet streamers form in the corona where magnetic fields of opposite polarity meet, and therefore a complex reconfiguration of the solar magnetic field is likely to occur due to magnetic reconnection. Magnetic reconnection can occur high up in the solar atmosphere (4–6 solar radii), and the collapse of newly formed magnetic loops can force the downward motion of coronal plasma. When densities are high enough, these plasma “inflows” are detected in coronagraphs in the vicinity of streamers and the coronal neutral line where the heliospheric current sheet (described next) forms (Wang et al., 2000). Multispacecraft studies using SoHO and STEREO images have recently shown that these inflows are associated with the release of density blobs in the slow solar wind (Sanchez‐Diaz et al., 2017).

Figure 1.4 Propagating brightness fluctuations (green/black) derived from COR‐2 observations during a dedicated deep‐field campaign. The fluctuations are present at all azimuths and times, with a wide range of brightnesses and lateral sizes. The fluctuations appear with smaller scales than the Sheeley blobs (discussed in the text); such a blob is observed off the north‐western limb in this image as the larger bright green feature.

(Source: Taken from DeForest et al., 2018.)

An important property of density blobs and plasmoids released in the solar wind is their multi‐scale and cyclic nature. Fourier analysis of brightness variations released from a highly tilted current sheet near solar maximum has shown that the large plasmoids are released from the Sun with characteristic time scales of about 19–20 hr (Sanchez‐Diaz et al., 2017). Similar spectral analysis has revealed characteristic 90 min timescales embedded within the larger plasmoids (Kepko et al., 2016; Viall et al., 2010; Viall & Vourlidas, 2015). DeForest et al. (2018) used deep‐field, high‐cadence coronagraph observations to show that there is still more substructure (shown in Figure 1.4)—both time dynamic on scales smaller than 90 min, and time stationary, filamentary streamer structures down to the resolution limit. The time‐dynamic characteristic scale sizes likely have different formation causes, such as inherent time scales due to the characteristics of coronal heating (Endeve et al., 2004), or waves (Pylaev et al., 2017).

Observations of the corona therefore show that highly dynamic streamers host transient processes associated with the likely continual emergence, redistribution, and removal of solar magnetic flux directly influencing the properties of the magnetic fields and particles of the solar wind (Owens et al., 2013; Sanchez‐Diaz et al., 2016). Magnetic reconnection occurring at the boundary of coronal holes was suggested to occur in the 1980s to explain the rigid rotation of coronal holes (Wang et al., 1988). Theoretical considerations (Fisk, 1996) and sophisticated numerical MHD simulations have highlighted the inevitable, complex evolution that takes place all along the corona’s open‐closed field boundaries (Antiochos et al., 2007; Linker et al., 2011; Lionello et al., 2005; Titov et al., 2009; Titov et al., 2011). The activity hosted by helmet streamers are the clearest example, with many different scales of plasma and magnetic field signatures generated by magnetic reconnection (Higginson et al., 2017; Higginson & Lynch, 2018; Lionello et al., 2005; Török et al., 2009).

Heliospheric imaging has also shown that not all structures at MHD scales measured in the solar wind form in the corona. DeForest, Matthaeus, Viall, and Cranmer (2016) showed turbulent density fluctuations setting in around 30 solar radii away from the Sun. These fluctuations coexist with outflowing helmet streamer plasmoids and are advected with the solar wind.

1.3. MEASUREMENTS OF THE SOLAR WIND IN THE INNER HELIOSPHERE

1.3.1. Bulk Properties and Large‐Scale Structures

The properties of the solar wind escaping the corona change throughout the solar cycle. This results from the evolving coronal magnetic topology that responds to the emergence and evolution of photospheric magnetic fields. This evolution alters both the magnetization and the bulk properties of the wind at different heliocentric latitudes. Spacecraft have measured in situ the properties of solar wind magnetic fields and particles at different locations in the heliosphere and over several decades. Figure 1.5 displays three dial plots showing the distribution of solar wind speeds with latitude measured by the Ulysses spacecraft during its three polar orbits. The sunspot number plotted below the dial plots is low in the left‐hand and right‐hand plots, marking the occurrence of solar minima (McComas et al., 2008). At these times, the Sun’s magnetic field is quasi‐dipolar, and the ambient solar wind is very clearly structured in at least two types of plasma flows. A fast wind is measured at high latitudes above coronal holes, and a more complex slow wind is measured at the low latitudes of solar streamers. As the solar cycle advances toward the activity maximum (middle panel), the polar coronal holes can disappear temporarily and the magnetic field evolves toward a non‐dipolar structure. In response to this process, the bulk speed loses the ordered latitudinal structure shown in Figure 1.5 (left). At solar maximum, the large‐scale spatial separation between the slow and fast solar wind is less clear as shown in the middle dial plot. Global numerical models of the solar wind show the clear link between these changes in the coronal magnetic topology and wind streams (Linker et al., 2011; Oran et al., 2013; Pinto et al., 2011; Pinto et al., 2016; van der Holst et al., 2014).

Figure 1.5 (a–c) Polar plots of the solar wind speed, colored by IMF polarity for Ulysses’ three polar orbits to indicate measured magnetic polarity. (d) Contemporaneous values for the smoothed sunspot number (black) and heliospheric current sheet tilt (red), lined up to match Figures 1.1a–c. In Figures 1.1a–c, the solar wind speed is plotted over characteristic solar images for solar minimum for cycle 22 (17 August 1996), solar maximum for cycle 23 (7 December 2000), and solar minimum for cycle 23 (28 March 2006). From the center out, we blend images from the Solar and Heliospheric Observatory (SoHO) Extreme ultraviolet Imaging Telescope (Fe XII at 1950 nm), the Mauna Loa K coronagraph (700–950 nm), and the SoHO C2 white‐light coronagraph.

(Source