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date: 22 September 2019

Solar Wind Origin

Summary and Keywords

The Sun continuously expels a fraction of its own mass in the form of a steadily accelerating outflow of ionized gas called the “solar wind.” The solar wind is the extension of the Sun’s hot (million-degree Kelvin) outer atmosphere that is visible during solar eclipses as the bright and wispy corona. In 1958, Eugene Parker theorized that a hot corona could not exist for very long without beginning to accelerate some of its gas into interplanetary space. After more than half a century, Parker’s idea of a gas-pressure-driven solar wind still is largely accepted, although many questions remain unanswered. Specifically, the physical processes that heat the corona have not yet been identified conclusively, and the importance of additional wind-acceleration mechanisms continue to be investigated. Variability in the solar wind also gives rise to a number of practical “space weather” effects on human life and technology, and there is still a need for more accurate forecasting. Fortunately, recent improvements in both observations (with telescopes and via direct sampling by space probes) and theory (with the help of ever more sophisticated computers) are leading to new generations of predictive and self-consistent simulations. Attempts to model the origin of the solar wind are also leading to new insights into long-standing mysteries about turbulent flows, magnetic reconnection, and kinetic wave-particle resonances.

Keywords: coronal heating, heliosphere, magnetic fields, solar atmosphere, solar corona, solar physics, solar wind, space weather


Solar Wind OriginClick to view larger

Figure 1. Perhaps the best way to “see” the solar wind is to block out the bright solar disk. This allows the much fainter solar corona, consisting of scattered light from charged particles in the million-degree outer atmosphere of the Sun, to become observable. This is a composite, processed illustration of the structures visible during the total solar eclipse of August 21, 2017 (adapted from original images obtained and processed by M. Druckmüller, P. Aniol, Z. Druckmüllerová, J. Hoderová, P. Štarha, and S. Habbal). Scattered light traces out magnetic field lines that have been stretched out by the expansion of the solar wind.

The solar wind is a tenuous gas surrounding the Sun that contains neutral atoms, positively charged ions, and free electrons. These particles accelerate away from the solar surface and fill the majority of the volume of the solar system, thus contributing to the fact that outer space is not a pure vacuum. The ever-expanding solar wind carries with it some of the Sun’s complex and multipolar magnetic field, which becomes stretched out along the mostly radial streamlines followed by the particles (see, e.g., figure 1). Far from the Sun, the typical outflow speeds of solar-wind particles range between about 250 and 800 km/s. The elemental composition of the solar wind is similar to that of the Sun’s interior: by mass, roughly 78% hydrogen, 20% helium, and 2% other elements—mostly oxygen, carbon, iron, neon, magnesium, nitrogen, and silicon. Collectively, these atoms, ions, and electrons compose a magnetized fluid that is an outer extension of the Sun’s atmosphere.

As it accelerates away from the Sun, the solar wind carries information about its origin in the solar corona. This information is encoded in the statistical distribution of particle speeds, which varies continuously as a function of time and radial distance away from the Sun. The bulk outflow velocity of the solar wind is the mean value of this distribution. There is always a nonzero random spread about the mean (i.e., the temperature) as well as asymmetries such as “skewness” that determine how momentum and energy are transported through the gas. Deep-space probes have measured these particle distributions in situ (i.e., locally, or in place) since the early 1960s, and they also provide a direct sampling of the electric and magnetic fields through which the particles travel. In addition, telescopes of various types observe solar photons that are scattered or emitted by the particles. These photons carry useful diagnostic information about the particle velocities and the magnetic field in the solar wind that is highly complementary to the in situ data.

The Sun loses approximately 1036 particles per second (i.e., about a billion kilograms per second) to the solar wind, which means that it is expected to lose roughly 0.02% of its total mass over its 10 billion year main-sequence lifetime. This is a small fraction, but the solar wind is believed to have been even more important to the Sun’s evolution in other ways. For example, the solar wind carries away some of the Sun’s angular momentum (see, e.g., Weber & Davis, 1967). Because of the magnetized nature of the fluid, the solar wind exerts a magnetic torque that is thought to have slowed the Sun’s rotation from an initial period of about 1 day to its current rotation period of about 26 days. Also, when the Sun was very young and still surrounded by a dense accretion disk, its rate of mass loss is thought to have been several orders of magnitude larger than it is now. The correspondingly denser “protosolar wind” may have been responsible for eroding away the primordial atmospheres of the inner planets (Lammer et al., 2012).

Additional motivations for studying the solar wind are often invoked from the standpoint of either pragmatism (i.e., space-weather forecasting) or as an accessible laboratory for fundamental physics research. On the pragmatic side, there is an ever-increasing need to understand how the Sun affects both technology and society. When the variable solar wind impacts the Earth’s magnetosphere, it can interrupt communications, disrupt power grids, damage satellites, and threaten the safety of humans in space (Feynman & Gabriel, 2000; Koskinen et al., 2017). Aside from this, the study of the physical processes at work in the solar wind has established a baseline of knowledge that assists in the understanding of the inner workings of more distant astrophysical systems. For example, theoretical insights about plasma heating mechanisms in the solar wind have been applied to other stars (Cranmer & Saar, 2011), the interstellar medium (Spangler, 1991), galaxy clusters (Parrish, McCourt, Quataert, & Sharma, 2012), and supermassive black holes (Sironi & Narayan, 2015).

Observational Overview

Historical Discovery

Since ancient times, it has been possible to observe both the launching of the solar wind (i.e., the corona during a total eclipse) and its eventual termination in the Earth’s atmosphere (i.e., the shimmering aurora borealis) with the unaided eye. From the late 19th century to the early 20th centuries, there arose considerable speculation that these phenomena were associated with one another. Carrington (1859) and others noted that the large solar flare of September 1, 1859, was soon followed by intense geomagnetic storms (rapid fluctuations in the Earth’s magnetic field) and disruptions along telegraph lines. Ennis (1878) speculated that the solar corona, the terrestrial aurora, and tails of comets were all similar manifestations of the same kind of “streaming forth of electricity.” Birkeland (1908), after collecting data for many years on polar expeditions, provided additional evidence that both auroral activity and geomagnetic storms “should be regarded as manifestations of an unknown cosmic agent of solar origin.” Serviss (1909) compared the “sheaves of light emanating from the poles” of the Sun’s corona to the field lines associated with the poles of a bar magnet. He also hypothesized that “there exists a direct solar influence not only upon the magnetism, but upon the weather of the earth.”

From the time of these earliest thoughts about a correlation between the Sun and the Earth’s local space environment, it took several decades to work out the precise means of causation. Chapman (1918) and others proposed that the Sun emits sporadic clouds or beams of charged particles into the vacuum of space. Hulburt (1937) proposed that bursts of ultraviolet radiation from the Sun were the primary means of exciting the aurora and geomagnetic storms. Pikel’ner (1950) realized that solar electrons can much more easily escape the Sun’s gravity than the heavier protons, and he speculated about the buildup of a net charge in the solar atmosphere. These ideas all had the benefit of being partially correct, but they did not yet fully come to grips with the idea of a plasma that is made up of both positively and negatively charged particles, and which behaves like a volume-filling gas or fluid. Biermann (1951) realized that the existence of cometary ion tails required there be a continuous outflow of “corpuscular radiation” (i.e., plasma) throughout the solar system.

At the same time as the above solar–terrestrial connections were being assembled, there was also evidence slowly building that the Sun’s corona is extremely hot. During the total solar eclipse of August 7, 1869, Harkness and Young first observed a bright green spectral line in the coronal emission at a wavelength of 530.3 nm. Because spectroscopists had not yet identified any chemical element that emits at this wavelength, Lockyer (1869) reported that these observations were “bizarre and puzzling to the last degree!” There was speculation that these observations implied the existence of a new element (“coronium”), and in subsequent eclipses a few dozen other mysterious coronal lines were discovered. The identification of these spectral features was a problem that persisted for at least half a century. Eventually, Grotrian (1939) and Edlén (1941, 1943) applied insights from the newly developed quantum theory to determine that these lines were coming from unusually high ionization states of iron, calcium, and nickel. These results confirmed that the corona must be a highly ionized plasma with a temperature of roughly 106 K. Alfvén (1941) assembled additional pieces of physical evidence that began to lead to modern ideas about coronal heating.

Parker (1958) brought together the idea of a hot corona with Biermann’s concept of a continuously outflowing collection of particles, and he found them to be inescapably linked. Parker realized that the high coronal temperature produced a strong outward force due to the gradient of gas pressure. This force is sufficient to counteract gravity and drive a time-steady accelerating flow of plasma that he named the solar wind. The initial model of Parker (1958) made the simplifying assumption of a constant coronal temperature, which is justifiable because the high thermal conductivity of a hot plasma leads to efficient spatial diffusion (i.e., smearing out) of temperature fluctuations (Chapman, 1957). Parker’s insightful calculation led to an implicit transcendental equation that contained both the radial outflow speed v and the radial distance r. Finding a closed-form analytic solution to this kind of equation (i.e., solving explicitly for v as a function of r) had to wait almost half a century for the development of new mathematical functions designed for these types of equations (see, e.g., Cranmer, 2004).

Parker’s solution for the solar wind was “transonic.” In other words, it involved a slow and subsonic expansion near the Sun, which begins to exceed the speed of sound at a critical radius in the corona, and then keeps accelerating as a supersonic flow at larger distances. This was not the only possible solution of Parker’s equations, and Chamberlain (1961) proposed that one of the other solutions—a solar breeze that always remains subsonic—was a more plausible one for nature to choose. There was also a controversy regarding Parker’s assumption that the plasma behaves as a continuous fluid. Models such as Chamberlain’s were based on a kinetic (i.e., collisionless particle-by-particle escape) approach that is still used in modeling planetary exospheres. At the time, these two approaches gave different answers, and it was not clear which one was correct.

Solar Wind OriginClick to view larger

Figure 2. Example solutions to Parker’s solar-wind equations: (a) a series of empirically constrained estimates for the temperature of the solar corona; and (b) the resulting radial dependence of outflow speed corresponding to each temperature curve. Radial distances are labelled in astronomical units (AU), where 1 AU = 215 solar radii. Parker’s critical radii are labelled with black circles. Note that, in the original theory, the full range of observed solar wind speeds (i.e., red, gold, and green curves) comes about only because of a relatively large range of coronal temperatures. However, the observed range of temperatures spans only the middle (gold) range of curves. Thus, additional physics is needed to explain the fastest (red) and slowest (green) solar-wind streams.

Prior to resolving these issues, Parker (1963) and others began to extend the original model to account for a radially varying coronal temperature. Depending on how and where the corona is heated, the high conductivity of the plasma allows for the temperature to increase to a maximum value in the corona, then decrease with increasing distance. Figure 2 illustrates this kind of model with an analytic temperature that obeys the dictates of its high conductivity both in the low corona (see, e.g., Rosner, Tucker, & Vaiana, 1978) and at large distances (Chapman, 1957). Hotter temperatures produce a critical radius closer to the Sun (i.e., the gas-pressure gradient starts overcoming gravity sooner) and a higher solar-wind speed in interplanetary space.

The Space Age: Verification and Exploration

It was fortuitous that Parker’s first theoretical model of the solar wind was published at the dawn of the Space Age, because the community had only a few years to wait for deep-space missions that ventured outside the Earth’s magnetosphere. The first in situ detection of solar wind particles came from a series of Russian Lunik and Venera probes between 1959 and 1961 (Gringauz, Bezruhkikh, Ozerov, & Rybchinskii, 1962). Also in 1961, the American satellite Explorer 10 measured plasma velocities and densities just outside the magnetopause. However, all of these measurements were rather brief. The ultimate confirmation of the existence of Parker’s solar wind—that is, the fact that it is always present and is always flowing at supersonic speeds—was provided by Mariner 2, which was sent to Venus in 1962 (Neugebauer & Snyder, 1966). Continuous data gathered over several months revealed the presence of alternating dense, low-speed (250–500 km/s) streams and tenuous, high-speed (500–900 km/s) streams.

Solar Wind OriginClick to view larger

Figure 3. (a) Looking down from the Sun’s north polar axis, magnetic field lines that fill the equatorial plane can be traced down to smaller source regions on the solar surface. The outer edge of this panel extends to a radial distance of 5 solar radii, or 0.02 AU. Field lines of positive (blue) and negative (gold) polarity were computed for the time period June–July 2007 by the Magnetohydrodynamics Around a Sphere (MAS) code (Linker et al., 1999); (b) As the wind accelerates out in the equatorial plane, fast and slow streams interact to form CIRs. Dark [light] colors illustrate low [high] density structures that follow the so-called “Parker spiral” field. The MAS model from panel (a) was used as input to a kinematic simulation that evolves the plasma from 0.02 to 1 AU (see, e.g., Riley & Lionello, 2011).

Soon after its discovery, additional large-scale structure in the interplanetary medium was detected. Wilcox and Ness (1965) analyzed data from IMP 1 (i.e., Explorer 18, which was renamed as the first Interplanetary Monitoring Platform) to reveal that the sign of the radial component of the magnetic field reverses its polarity every few days. This pattern of oppositely directed “magnetic sectors” also tends to recur roughly every 27 days, which is close to the Sun’s equatorial rotation period. Figure 3a illustrates the current understanding of these sectors; they each map back to discrete source-regions in the solar corona of the same polarity. The high conductivity of plasma in the corona and solar wind ensures that the particles and the lines of magnetic force remain closely tied to one another, so that the entire system rotates together. Alfvén (1957) and Parker (1958) realized that this effect would lead to radially flowing streams of particles winding up the magnetic field lines into rotating Archimedean spirals. This is analogous to the way that a spinning lawn-sprinkler emits radial jets of water that are observed to be twisted into spiral-shaped streaklines. Slower wind streams are more tightly wound than faster streams. This means that when fast and slow streams are emitted at neighboring longitudes, they eventually interact with one another to form co-rotating interaction regions (CIRs). Figure 3b shows that high-density regions (i.e., compressions) occur when fast streams catch up with slow streams, and that low-density regions (rarefactions) occur when the slow streams lag behind.

Spatial and temporal variability in the solar wind extends down to very small scales. Belcher and Davis (1971) discovered that so-called Alfvén waves (i.e., transverse oscillations in the magnetic field that do not disturb the density or pressure like sound waves do) are ubiquitous in the solar wind. Non-linear interactions between different types of fluctuations create a chaotic “turbulent cascade” in which large eddies are broken up into ever-smaller eddies (see, e.g., Bruno & Carbone, 2013). In fact, the eddies keep cascading over at least eight orders of magnitude in size until they find themselves at the scales on which individual charged particles interact with one another and with the magnetic field. In regions of low density—where inter-particle collisions become infrequent—these effects manifest themselves as departures from thermal equilibrium (i.e., individual particle species drifting past one another at different speeds, and temperatures measured along the magnetic field being unequal to temperatures measured transverse to the field; see, e.g., Marsch, 2006).

In the decades since the solar wind’s discovery, there have been many other deep-space missions that have studied it in more detail. The earliest missions tended to stay close to the Earth’s orbit at a heliocentric distance of 1 astronomical unit (AU). The twin Helios probes plunged inside the orbit of Mercury to 0.28 AU (Marsch, 1991). The Voyager probes are now past 120 AU—approximately three times more distant than Pluto—and in 2019 are still sending back data (Richardson et al., 2017). In the 1990s, Ulysses became the first spacecraft to pivot away from the solar system’s ecliptic plane and measure the solar wind over the north and south poles of the Sun (Marsden, 2001). The Ulysses, SOHO (Solar and Heliospheric Observatory), Wind, and ACE (Advanced Composition Explorer) missions contained instruments that were able to measure the precise abundances and ratios of ionization charge states for multiple elements heavier than hydrogen and helium (e.g., Zurbuchen, 2007). Because these “composition” signatures do not evolve as rapidly as other properties of the plasma (e.g., temperature, total density, and flow speed), they are useful tracers of the coronal origins of solar-wind streams.

A survey of solar-wind measurements would not be complete without also including telescopic remote-sensing. With Lyot’s development of the coronagraph in the 1930s, it became possible to observe scattered light from the solar corona at times other than total solar eclipses (see, e.g., Billings, 1966). Coronagraphic measurements in broadband visible light are sensitive to the number of free electrons along various lines-of-sight through the corona. Thus, it has become possible to use techniques derived from medical tomography to determine the three-dimensional distribution of electron density (Frazin, Lamy, Llebaria, & Vásquez, 2010). Rapid-cadence sequences of coronagraphic images have also revealed the presence of low-level density fluctuations that drift to larger radii as a function of time. These structures are believed to act like passive “leaves in the wind” and thus probe the acceleration profile of the flow (Sheeley et al., 1997; Cho et al., 2018). Lastly, there are also coronagraphic instruments that disperse the incoming light into a spectrum. This allows hundreds of bright coronal emission lines (including those that were initially mistaken for “coronium”) to be used as diagnostics of elemental abundances, plasma-velocity distributions, and turbulent fluctuations. In the regions of the corona undergoing the fewest inter-particle collisions, the observations show similar departures from thermal equilibrium as are seen by the in situ instruments (see, e.g., Kohl, Noci, Cranmer, & Raymond, 2006).

Theoretical Understanding

What Heats the Solar Corona?

The source of solar-wind acceleration proposed by Parker (1958, 1963)—that is, an outwardly directed gas-pressure gradient force due to the existence of the hot corona—is still believed to be responsible for much of what is observed. Early questions about (1) the viability of the transonic solution to the conservation equations, and (2) mutual consistency between fluid-based and kinetic approaches to modeling the flow, were resolved more or less in Parker’s favor (Velli, 1994; Echim, Lemaire, & Lie-Svendsen, 2011). Thus, the most substantial remaining problem is how to produce the required million-degree coronal temperatures.

The ultimate source of energy for heating the solar corona is the convection zone that churns just below the surface. Only about 1% of the kinetic energy in the rising and falling motions of convective “granules” needs to be delivered up to the corona and subsequently converted into heat. Because more highly magnetized regions of the Sun appear to receive more coronal heating, it is widely believed that the upward delivery of energy must involve temporary storage in the form of magnetic energy. Thus, a reasonable way of expressing the maximum available energy is to evaluate the upward component of the Poynting flux (i.e., the energy transfer per unit area, per unit time, by a magnetic field in motion). An upward Poynting flux at the solar surface can be achieved by either (1) transverse jostling of mostly vertical field lines, or (2) upward flow of mostly horizontal field lines. Of course, for nearly all regions on the actual Sun, the field is so twisted and tangled—and the underlying motions are so complex—that the Poynting flux probably has nontrivial contributions from both effects.

The irreversible conversion of magnetic energy to heat is most efficient when there arise structures on very small spatial scales. In other words, an increase in the random thermal motions of particles is most likely to occur when there is activity on scales commensurate with the (microscopic) random inter-particle collisions. For the corona, there is still no agreement on how this happens; see recent reviews by Parnell and De Moortel (2012), Schmelz and Winebarger (2015), and Cranmer, Gibson, and Riley (2017). There are two overall schools of thought that depend on the relative values of two key time-scales. First, there is the so-called Alfvén travel-time tA, which is the time that it takes a small wavelike fluctuation to traverse a significant distance along the coronal magnetic field. Second, there is the photospheric footpoint time-scale tph, which is a characteristic time over which the convective motions can make major changes to the field at the bottom of the corona. Thus, there are two limiting cases:

  1. 1. If tph > tA, then the corona has ample time to respond to the slow underlying motions. The coronal magnetic field relaxes into a succession of increasingly braided and stressed configurations. At any arbitrary time in such a system, there will be a selection of locations at which the magnetic complexity exceeds a stability threshold. Each of these locations will then undergo a rapid, explosive burst of current dissipation known as magnetic reconnection (Parker, 1972, 1988). These bursts of heating are often called “nanoflares,” and the models that describe these kinds of quasi-static equilibria are called direct-current (DC) heating theories. These mechanisms are invoked more often for closed magnetic loops than for the open field lines that feed the solar wind (see, e.g., the “flux-tube tectonics” model of Priest, Heyvaerts, & Title, 2002).

  2. 2. On the other hand, if tph < tA, then the rapid footpoint driving creates fluctuations that propagate along the coronal magnetic field lines in the form of waves, shocks, and turbulent eddies (see, e.g., Osterbrock, 1961; Roberts, 2000). Small spatial scales can be generated when the waves interact with background gradients in density (e.g., reflection or refraction at sharp boundaries) or when neighboring waves travel at different speeds through an inhomogeneous plasma (e.g., “phase mixing”; see Heyvaerts & Priest, 1983). Eventually, the waves damp out, either via particle–particle collision effects (i.e., viscosity, thermal conductivity, or electrical resistivity) or via wave-particle interactions that occur when there are departures from thermal equilibrium. The damped wave energy is converted to heat, and the models that utilize this source of energy are often called alternating-current (AC) heating theories.

One source of ongoing difficulty with the above ideas is that the observed time-scales often take on different values in different regions. Statistically, both tph and tA tend to have broad distributions that overlap with one another, so it is unclear whether there are very many strictly AC or DC regions. Also, some observational clues once believed to be unique to one paradigm seem to have the characteristics of the opposite paradigm. For example, high-resolution ultraviolet imaging has seen signs of both the bursty nanoflare events and tightly braided magnetic fields characteristic of DC theories. However, these tend to coincide with rapid and fluctuating flows that are expected from AC theories (see, e.g., Cirtain et al., 2013; Testa et al., 2013). This kind of coronal activity can be described using the unifying language of a turbulent cascade (see, e.g., Heyvaerts & Priest, 1984; Milano, Gómez, & Martens, 1997; van Ballegooijen, Asgari-Targhi, & Beger, 2014; Velli, Pucci, Rappazzo, & Tenerani, 2015).

What Else Accelerates the Wind?

Figure 2 shows that there are high-speed wind streams (i.e., exceeding 600 km/s) that are difficult to explain with Parker’s original gas-pressure driving theory alone. These fast streams tend to contain relatively high fluxes of Alfvén waves that propagate away from the Sun. Thus, Belcher (1971) and Alazraki and Couturier (1971) proposed that these oscillations may exert an additional outward-directed wave pressure that provides extra acceleration to the wind. Several years earlier, Bretherton and Garrett (1968) showed that waves propagating through an inhomogeneous medium can exert a nondissipative net force on the fluid. This force can be expressed as the gradient of a pressure-like quantity; the effect is somewhat analogous to the way that electromagnetic waves carry momentum and exert pressure on matter. Recent models of the solar wind that include wave pressure have been successful in predicting the properties of high-speed solar-wind streams (see, e.g., Cranmer, van Ballegooijen, & Edgar, 2007; Ofman, 2010; Oran et al., 2013).

In the lowest-density parts of the corona—in which inter-particle collisions are rare—departures from thermal equilibrium can provide additional ways of accelerating the solar wind. There can be “two-temperature” distributions of particle velocities, in which the spread of random velocities perpendicular to the magnetic field lines is larger than the spread parallel to the field lines (Marsch, 2006). In this case, particles that flow away from the Sun undergo a similar magnetic mirror effect as particles in the Earth’s Van Allen radiation belts. As individual particles travel from strong to weak magnetic-field regions, their perpendicular gyrations become redirected to be parallel to the (mostly radially oriented) field lines. This provides extra kinetic energy for the outflowing solar wind.

The slowest streams of the solar wind are seen to flow away from the Sun at speeds down to about 250 km/s. The original Parker (1958) theory would have demanded coronal temperatures of only about 500,000 K to create a wind this slow, but such low temperatures do not tend to be observed. It is possible that these regions are the result of energy loss in the stream–stream collisions that form co-rotating interaction regions (CIRs). Alternately, they may be associated with highly distorted bundles of magnetic field lines, that is, far from the radially oriented rays assumed in Parker’s spherically symmetric model. These geometrical effects are described in more detail in the section titled “Forecasting the Wind Speed.” Lastly, Kasper, Stevens, Lazarus, Steinberg, and Ogilvie (2007) discovered an intriguing fact about the elemental composition of the slowest solar wind: the relative abundance of helium decreases rapidly as one approaches the low-speed limit of 250 km/s from above. No solar wind at all is observed with speeds below the value at which the helium abundance would be extrapolated to zero. Thus, there is speculation that the presence of helium is a crucial piece of the solar wind acceleration puzzle.

What Determines the Sun’s Mass Loss?

Parker’s (1958) theory, and its many extensions, are largely successful in predicting the acceleration and speed of the solar wind. However, the theory does so without putting any clear limits on how many particles are ejected per unit time. Observations and models have shown that the full Sun emits roughly 2 × 10−14 solar masses per year in the solar wind at the minimum of the sunspot activity cycle, and roughly 3 × 10−14 solar masses per year at the maximum of the cycle (Wang, 1998). Note that the Sun’s interior mass is also depleted at a slightly higher rate (6.7 × 10−14 solar masses per year) by the conversion of hydrogen to helium via nuclear fusion. However, this latter rate of mass depletion is observable only in the form of the solar luminosity (mass–energy release as photons) and not as a component of the solar wind.

Given that both the solar magnetic field and the total rate of coronal heating vary by orders of magnitude over the 11-year activity cycle, it is somewhat surprising that the mass-loss rate varies by only about 50%. However, the energy deposited as coronal heating is very efficiently transported away from where it originates, both inward (in the form of thermal conduction) and outward (as the solar wind). These effects essentially smear out the effects of any large changes in the heating. In other words, this acts as a kind of “thermostat” that limits the variability of the coronal temperature and pressure (see, e.g., Rosner et al., 1978; Hammer, 1982; Leer, Holzer, & Flå, 1982). The thermostat effect is a consequence of energy conservation, and the requirement for a steady-state balance between gains (heating) and losses (both transport and actual cooling via the emission of photons) is what sets plasma density—and thus also the mass-loss rate—to vary as weakly as it does.

Although the idea of thermal energy balance has been the dominant explanation for the Sun’s mass loss, another idea has been suggested. High-resolution observations of the solar chromosphere (a relatively cool layer between the photosphere and corona) have revealed the presence of a wide range of narrow, short-lived features known variously as spicules, jets, fibrils, and mottles. Many of these structures show a rapid upward surge of plasma, followed by much (or all) of that plasma falling back down. It has been suggested that some of this plasma continues traveling up and is eventually heated to become the corona and the base of the solar wind (see, e.g., Moore, Sterling, Cirtain, & Falconer, 2011; McIntosh, 2012). This is related to the idea of interchange reconnection in the low corona, in which magnetic loops emerge from below the solar surface, undergo magnetic reconnection with neighboring regions, and drive jet-like surges of plasma up into the corona and solar wind (Fisk, Schwadron, & Zurbuchen, 1999; Yang et al., 2013). Although coronal jets are indeed observed, they appear as bright and narrow features in images because they occupy only a small fraction of the coronal volume at any time. Thus, it remains uncertain whether they can be responsible for the majority of the plasma mass of the corona and solar wind.

Forecasting the Wind Speed

As summarized in the “Introduction,” there is a definite societal need to for predictions of the future state of the solar wind. Putting aside violent events such as solar flares and coronal mass ejections (CMEs), it is also important to know how the Sun produces its “ambient” distribution of fast and slow streams. In turn, this requires knowing which magnetic structures in the corona connect to which kinds of wind streams at 1 AU. The highest-speed streams appear to be connected to the largest coronal holes on the Sun’s surface (Cranmer, 2009). These are the darkest and least active regions of the Sun, but they expand out majestically over the north and south poles over much of the solar cycle (see figure 1). On the other hand, the slow solar wind is associated with a range of observed magnetic structures: for example, cusp-like streamers often seen at low latitudes, fan-like active regions associated with sunspots, quiet regions with only intermittent open fields, and the outer edges of coronal holes (see, e.g., Abbo et al., 2016). In addition, Antiochos, Mikić, Titov, Lionello, and Linker (2011) highlighted the possible importance of a web-like network of magnetic separatrix surfaces that appear to extend up into the source regions of the slow solar wind. Because these coronal structures all undergo substantial time variability, their precise relative contributions to the slow wind are still not known.

No matter the dominant magnetic sites of origin for solar-wind streams, it is now well known that there is a relationship between the speed of the wind and the flux-tube geometry of the field-lines that connect that stream down to the solar surface. Levine, Altschuler, and Harvey (1977) and Wang and Sheeley (1990) independently discovered that the wind speed at 1 AU is correlated inversely with the amount of transverse flux-tube expansion between the surface and a reference point in the mid-corona (see also Arge & Pizzo, 2000). Bundles of field lines that trumpet out abruptly from a small source region on the Sun generally produce slow wind. Bundles of field lines that undergo only moderate lateral expansion (such as the central regions of coronal holes) generally produce fast wind. The physical origin of this effect is suspected to be related to the emission of Alfvén waves at the base of the corona (Kovalenko, 1981; Wang & Sheeley, 1991). If every region on the surface produces the same energy flux (i.e., power per unit area), then regions that expand greatly—from small patches on the surface to large areas in the outer corona—will receive lower total amounts of wave energy than the regions that map down to larger patches on the surface and do not expand as much. Thus, the low-expansion regions with the highest amounts of coronal wave energy would then be heated more vigorously (from, for example, the turbulent cascade of those waves), and thus be subject to higher degrees of solar-wind acceleration. The relationship between flux-tube expansion and solar-wind speed forms the backbone of much practical space-weather forecasting (see, e.g., Riley, Linker, & Arge, 2015).

Conclusions and Future Prospects

Although there have been substantial accomplishments in both observations and theory since the late twentieth century, there is much more to be done in order to reach a complete understanding of the origin of the solar wind. Advances in computational efficiency have allowed the community to build increasingly sophisticated numerical simulations of the corona and solar wind (e.g., Gressl et al., 2014; Carlsson, Hansteen, Gudiksen, Leenaarts, & De Pontieu, 2016; Chen, Rempel, & Fan, 2017; Yalim, Pogorelov, & Liu, 2017; Gombosi, van der Holst, Manchester, & Sokolov, 2018; Shoda, Yokoyama, & Suzuki, 2018), and the use of these kinds of models in real-time forecasting is on the horizon. Multidimensional simulations are increasingly making use of data assimilation (i.e., constant updating based on the most recent observations) to improve their effectiveness and their fidelity to nature.

In addition to theoretical advances, there are also new vistas in observational capability coming soon. The Parker Solar Probe (Fox et al., 2016) is expected to revolutionize the community’s conception of the inner heliosphere by performing in situ sampling closer to the Sun than any other space mission. The European Space Agency’s (ESA) Solar Orbiter mission (Müller, Marsden, St. Cyr, & Gilbert, 2013) will provide new access to both the solar poles and the inner heliosphere with combined in situ and remote-sensing instruments. The Daniel K. Inouye Solar Telescope (DKIST; Tritschler et al., 2016) will be the largest-diameter solar telescope in history and it will measure the solar corona’s magnetic field with unprecedented sensitivity. There are also plans in development for placing solar-wind monitors at various locations in the solar system—for instance, at the Earth–Sun L5 point or over the solar poles—that would help fill in huge observational gaps and thus lead to more accurate space-weather forecasting.

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