1 Introduction

Earth’s magnetic field carves out a cavity in the oncoming solar wind known as the magnetosphere. Because the magnetosphere extracts all of the mass, momentum, and energy that powers geomagnetic storms from the solar wind, quantifying and understanding the flow of these quantities from the Sun outward through the heliosphere, through the Earth’s magnetosphere, and into the Earth’s ionosphere is one of the primary goals of the heliophysics discipline. Similar objectives govern the planetary discipline which seeks, amongst other tasks, to determine the nature of the solar wind’s interaction with comets and the other planets within our solar system, and in particular to quantify the role that plasma processes play in the loss of their atmospheres. Once the conditions governing the occurrence patterns of the various fundamental processes (including reconnection, diffusion, instabilities, particle acceleration, and ion-neutral interactions) that control the mass, energy, and momentum flow are well understood, it will become possible to construct numerical simulations that provide accurate space weather predictions for the immediate environment of the Earth and other solar system objects (e.g., Bertucci et al. 2011).

Figure 1 presents results from state-of-the-art hybrid code simulations for the plasma interactions that occur in the vicinity of Venus, Mars, and the Earth. From a global perspective, the density structures, and thus the processes that govern these interactions exhibit many similarities. A bow shock (BS) separates the higher density magnetosheath plasma of solar wind origin from the solar wind itself. A sharp magnetopause or ionopause (I) separates the magnetosheath from the planetary obstacle, whether it be the high density ionospheres with plasmas of planetary origin at Venus and Mars or the low density magnetosphere at Earth. The panels for Venus and Mars show boundary locations for a stable interplanetary magnetic field (IMF) transverse to the Sun-planet line, while the third panel shows boundaries near Earth during the passage of a solar wind tangential discontinuity (TD).

Fig. 1
figure 1

Cuts through hybrid simulations of the solar wind-magnetosphere interaction showing the density of plasma with a solar wind origin at Venus (upper left panel, Bößwetter et al. 2007), Mars (upper right panel, Shimazu 2001b), and Earth (lower panel, Omidi and Sibeck 2007). The panels for Venus and Mars show boundary locations for a stable IMF transverse to the Sun-planet line, while the third panel shows boundaries near Earth during the passage of a solar wind TD at which the IMF rotates from northward and antisunward to dawnward. Here BS stands for bow shock and I for ionopause. Distances in the second panel are measured in planetary radii, in the first and third panel they are measured in terms of the ion skin depth (\(c/\omega_{pi} = c [4{\pi} n_{p} e^{2}/M]^{-{1\over 2}}\sim100~\mbox{km}\) for \(n=10~\mbox{cm}^{-3}\), where \(c\) is the speed of light, \(M\) the mass of a proton, \(n_{p}\) the proton density, and \(e\) the charge of an electron). Densities in the first and third panel have been normalized to those in the solar wind. Note the multiple shock structures at Venus, the north/south asymmetries in bow shock and ionopause locations at Mars, and the complex shock structure at Earth

Many micro- to macro-scale processes have been predicted and observed to occur in the vicinity of the bow shock and magnetopause, as well as throughout the foreshock, magnetosheath, and outer magnetosphere. These processes are often identified on the basis of the diagnostic density structures that they generate. Macroscale structures include the bow shocks, magnetosheaths, and either the ionopauses or magnetopauses that stand upstream from both comets and planets. The location and motion of these boundaries depend not only on the time-varying conditions within the solar wind but also on conditions within the magnetospheres and ionospheres. Mesoscale features include dawn/dusk asymmetries in foreshock and magnetosheath parameters, waves and riplets driven by variations in solar wind parameters or instabilities on the boundaries, boundary layers of intermingled magnetosheath and magnetospheric or ionospheric plasma, and cusps filled with magnetosheath-like plasma that link Earth’s magnetopause to its ionosphere and atmosphere. Microscale features include the kinetic structures generated by wave-particle interactions within the foreshock and the structure of the bow shock and density variations associated with magnetosheath waves.

The significance of each interaction process depends upon its spatial extent and the solar wind/magnetospheric conditions under which it occurs. While statistical studies of in situ observations can provide considerable information concerning the occurrence patterns of various phenomena, reconstructing the global configuration of density structures from isolated in situ measurements is no easy task. Global magnetohydrodynamic and, more recently, hybrid kinetic simulations provide considerable insight, but need validation by equally global measurements.

Pending the launch of constellation-type missions with thirty or more spacecraft in a wide array of orbits (e.g., The Magnetospheric Constellation, MC: Global Dynamics of the Structured Magnetotail, NASA 2004), imaging affords the best (and certainly the most cost-effective) means of (1) determining the overall configuration of the Earth’s magnetosphere, (2) identifying the extent and significance of the processes governing the solar wind-magnetosphere interaction on the basis of their diagnostic plasma density signatures, and (3) validating the numerical simulations. Missions like DE-1 (Frank et al. 1981), Viking (Anger et al. 1987), Freja (Murphree et al. 1994), Polar (Frank et al. 1995; Imhof et al. 1995; Torr et al. 1995), and IMAGE (Mende et al. 2003) employed visible, ultraviolet, and X-ray imagers to take global pictures of the auroral oval, a region to which many of the most basic processes in the magnetosphere map. However, it can be difficult to determine both the nature of the processes and the locations of distinctive features in the magnetosphere that map to features in the auroral oval. The need for global images of the magnetosphere led to the launch of IMAGE and TWINS. These missions took extraordinarily fascinating and instructive images of the plasmasphere in extreme ultraviolet, of the cusp and subsolar magnetosheath in low-energy neutral atoms, of the auroral oval in previously unobserved far ultraviolet wavelengths, and of the ring current in higher energy neutral atoms. Discoveries included plasmaspheric shoulders and notches (Darrouzet et al. 2009), surprisingly slow plasmaspheric rotation (Burch et al. 2004), a hot oxygen geocorona (Wilson et al. 2003), and persistent proton auroras (Frey et al. 2003).

Observations of the global solar wind-magnetosphere interaction suitable for direct comparison with the predictions of global numerical models are now within reach. Operating from vantage points up to \(49~\mbox{R}_{E}\) from Earth, the IBEX-Hi imager (Funsten et al. 2009) on the spinning (\(\sim4~\mbox{rpm}\)) Interstellar Boundary Explorer spacecraft (IBEX, McComas et al. 2009) has returned rastered images of the bow shock, magnetopause, and cusps in 0.9–1.5 keV energetic neutral atoms (ENAs), primarily hydrogen. The solar wind protons acquire electrons from exospheric hydrogen atoms and then proceed in their pre-exchange directions. Because the decelerated and thermalized solar wind protons gyrate around magnetosheath magnetic field lines, the pre-existing directions are effectively random over the expected scale lengths of magnetosheath phenomena, and the ENA flux seen in any direction is approximately proportional to the integrated line-of-sight (LOS) product of the plasma ion and exospheric neutral densities. Figure 2 presents examples for the magnetosheath (Fuselier et al. 2010) and cusp (Petrinec et al. 2011).

Fig. 2
figure 2

ENA images of the dayside magnetosphere from the IBEX mission. The left panel presents measurements of ENAs from the subsolar magnetosheath (adapted from Fuselier et al. 2010), while the right panel shows ENAs from the cusps (Petrinec et al. 2011)

Strikingly different ENA flux levels are observed on LOS integrations that (1) remain solely in the low plasma and low neutral density solar wind, that (2) pass through the high plasma and moderate neutral density magnetosheath, that (3) pass through the high plasma and high neutral density cusps, and (4) that pass through the very low plasma and high neutral density equatorial or polar magnetosphere. Furthermore, the energies, composition, flux, and direction of the ENAs arriving at the observing location provide important information concerning the processes occurring at remote magnetospheric locations (Taguchi et al. 2004; Collier et al. 2005a; Hosokawa et al. 2008).

On the other hand, the \(7^{\circ}\times7^{\circ}\) single pixel IBEX-Hi imager requires times ranging from 11 to 20 hours to raster individual global ENA images, with inherent spatial resolutions in the noon-midnight meridional plane ranging from \(3.7~\mbox{R}_{E}\) for spacecraft locations just outside the bow shock to \(6.1~\mbox{R}_{E}\) at \(49~\mbox{R}_{E}\) apogee. By contrast, cadences on the order of minutes to tens of minutes and spatial resolutions less than \(1~\mbox{R}_{E}\) are needed to capture the dynamics of the processes that govern the solar wind-magnetosphere interaction at the bow shock and magnetopause. Even if instantaneous global snapshots could be taken, the finite times-of-flight required for individual ENAs to arrive at the observing instrument would result in individual images representing the convolution of particles with different energies coming from different locations at different times.

An alternative method for imaging the magnetosphere offers the potential to obviate these problems. Exospheric neutral charge exchange with high charge state solar wind ions generates soft X-rays with energies from 0.05–2.0 keV. Currently existing wide field-of-view (FOV) soft X-ray telescopes provide an opportunity to image not only the dayside solar wind-terrestrial magnetosphere interaction, but also the interactions that occur at the Moon, Venus, Mars, and comets. This paper begins with a review of those scientific topics raised by modeling and past in situ missions that can be addressed by imaging missions. It then discusses the physical processes governing the generation of soft X-rays, in particular charge exchange with high charge state solar wind ions. Numerical simulations employ models for the solar wind composition, exosphere, solar wind-magnetosphere interaction, and soft X-ray background to predict the integrated LOS emission intensities observable by wide FOV soft X-ray imagers and define the cadence and spatial resolution required from such an imager. A review of previously reported observations by narrow FOV astrophysical telescopes demonstrates that the emissions are present at the predicted level from all of the proposed targets. Wide FOV soft X-ray telescopes capable of making global observations with the required spatial resolution and cadences have already flown and are scheduled for forthcoming missions. The features seen within the global images can be readily associated with density structures observed by in situ spacecraft on suitable orbits. The paper concludes with comments concerning prospects for wide FOV soft X-ray telescopes.

2 Scientific Objectives

Global images of the soft X-rays generated when high charge state solar wind ions (e.g., \(\mbox{C}^{6+}\), \(\mbox{O}^{7+}\), \(\mbox{O}^{8+}\), \(\mbox{Fe}^{12+}\)) exchange charges with neutrals (e.g., H, H2O) can provide crucial information concerning the nature of the solar wind’s interaction with planetary atmospheres and magnetospheres, including those of the Earth, Venus, Mars, the Moon, and comets. As illustrated in Fig. 3, the reason for this is that the processes governing the interaction of the solar wind with these heliospheric obstacles generate a host of plasma density structures that can be used to diagnose the nature of those interactions. At the Earth, the size, shape, structure, and motion of the magnetopause and cusps provide important information concerning the global characteristics of magnetic reconnection, the strength of various magnetospheric current systems, and the response of the magnetosphere to varying solar wind and foreshock input. Observations of transients at the magnetopause and in the cusps quantify their extent and occurrence patterns, hence their significance to the overall interaction. Observations of the magnetosheath structure and its time variability provide the outer boundary conditions for the magnetosphere. The location of the bow shock yields information concerning the thermodynamics of the collisionless solar wind, while the structure of the bow shock defines its ability to reflect and energize particles, a fundamental heliospheric process. Observations of the foreshock are needed to understand and quantify the effects of the particles accelerated at the bow shock upon the bulk parameters of the incoming solar wind and therefore upon the overall solar wind-magnetosphere interaction.

Fig. 3
figure 3

Plasma structures generated by the solar wind’s interaction with Earth’s magnetosphere: solar wind (SW), bow shock (Bshock), and magnetopause (MP). Adapted from Wiltberger et al. (2015)

There are parallel research problems to be addressed by imaging comets, the Moon, Venus, and Mars. These topics concern the interaction of the solar wind with obstacles that have little or no intrinsic magnetic field. In the cases of comets, Venus, and Mars, studies that focus on the location, structure, and motion of the bow shock and ionopause yield information concerning atmospheric loss rates. In the case of the Moon, studies focus upon the structure, composition, and sources of the tenuous lunar exosphere. This section describes potential research questions.

2.1 The Earth

We begin by considering those questions concerning the Earth’s magnetopause, cusps, transients at the magnetopause and in the cusps, the magnetosheath, bow shock, and foreshock that can be diagnosed with the help of information concerning plasma density structures deduced from soft X-ray observations. We then address those questions concerning the processes that occur at comets, Venus, Mars, and the moon that can also be answered with the help of soft X-ray images.

2.1.1 Earth’s Magnetopause

A host of factors, including the solar wind thermal and dynamic pressures, the IMF latitude and cone angle, the dipole tilt, and the strength of various current systems within and bounding the magnetosphere determine the location of the magnetopause. Although they can predict widely divergent magnetopause locations for the same solar wind conditions (Samsonov et al. 2016), global magnetohydrodynamic simulations all indicate that the solar wind dynamic pressure and north/south component of the IMF are the most important factors determining magnetopause location (Lu et al. 2011). As illustrated in the top panel of Fig. 4, empirical studies based on large numbers of magnetopause crossings paired with time-averaged solar wind measurements suggest that the magnetopause expands and contracts in a self-similar manner in response to variations in the solar wind dynamic pressure (Sibeck et al. 1991; Roelof and Sibeck 1993; Lin et al. 2010; Wang et al. 2013). Some case studies disagree (Stüdemann et al. 1986). Both case studies (e.g., Kaufmann and Konradi 1969) and numerical simulations (Samsonov et al. 2015) confirm that the response of the magnetopause to step-function changes in the solar wind pressure is more complicated than self-similar contractions and expansions.

Fig. 4
figure 4

Results from an empirical model for the locations of the equatorial magnetopause as a function of (upper panel) 5 solar wind pressures (0.54–0.87, 0.87–1.47, 1.47–2.60, 2.60–4.90, and 4.90–9.90 nPa) and (lower panel) 6 values of IMF Bz (−6 to −4, −4 to −2, −2 to 0, 0 to 2, 2 to 4, and 4 to 6 nT) (Sibeck et al. 1991). The plots are in geocentric solar ecliptic (GSE) coordinates in Earth radii (\(\mbox{R}_{E}\)) with \(\mbox{R}=\sqrt{y^{2}+z^{2}}\)

By contrast, in response to changes in the IMF orientation, the dayside magnetopause moves Earthward (Aubry et al. 1970), the cusps move equatorward (Newell et al. 1989), and the magnetotail flanks move outward (Maezawa 1975) during intervals of southward IMF orientation, thereby producing a blunter magnetosphere with a greater magnetopause flaring angle. This erosion, or inward motion of the dayside magnetopause and outward motion of the magnetotail magnetopause, can be attributed to magnetic reconnection, a process that removes magnetic flux from the dayside magnetosphere and adds it to the magnetotail, although it has recently been noted that a (small) portion of the inward motion may result from the enhancements of the pressure near the subsolar magnetosheath known to occur for the blunter magnetopause shapes during intervals of southward IMF orientation (Shue et al. 2013). Wiltberger et al. (2003) propose that magnetopause erosion results from (rather than causes) enhanced cross-tail currents. Soft X-ray images will provide an opportunity to determine the instantaneous shape of the global magnetopause and define its evolving response to solar wind variations, thereby distinguishing between the possibilities outlined above.

Because it is the dominant process enabling the transfer of solar wind mass, energy, and momentum into the magnetosphere, understanding reconnection is a fundamental heliophysics objective. In conjunction with solar wind observations, magnetopause locations and shapes can be used to deduce the magnetic field strengths just inside the magnetopause, the strengths of the relevant magnetospheric current systems, the amount of flux eroded from the dayside magnetosphere, and as a result, the global response of reconnection to varying solar wind conditions (Sibeck et al. 1991). At rest, the magnetopause lies along the locus of points where the sum of thermal and magnetic pressures balance in the magnetosheath and magnetosphere. Under both elastic and inelastic reflection hypotheses, the magnetosheath pressure applied locally to the dayside magnetopause is proportional to the fraction of the solar wind dynamic incident upon the flaring magnetopause surface (e.g., Spreiter et al. 1966). With the exception of the cusps, where plasma pressures are high, the total pressure applied by the magnetosheath to the magnetopause is balanced almost exclusively by the magnetic pressure just inside the magnetopause. However, the magnetic fields that contribute to this magnetic pressure are themselves just the sum of contributions from all magnetospheric current systems.

Thus, together with a measure of the solar wind dynamic pressure, soft X-ray observations of the location and shape of the magnetopause can be used to infer magnetic field strengths just inside the magnetopause and in turn variations in magnetospheric current systems as a function of solar wind conditions. In the case of reconnection, the relevant current systems are the Region 1 Birkeland current and, to a much lesser degree, the cross-tail current systems (Maltsev and Lyatsky 1975; Tsyganenko and Sibeck 1994). Operating in tandem, these current systems reduce dayside magnetospheric magnetic field strengths, transfer magnetic flux to the magnetotail, and allow the dayside magnetopause to move inward during intervals of southward IMF orientation. With their strengths inferred from observations of the dayside magnetopause location, the amount of flux eroded by reconnection from the dayside magnetosphere can be determined for any combination of solar wind or geomagnetic parameters (Sibeck et al. 1991; Shue et al. 2001).

Observations of magnetopause motion can be used to determine the time-dependence of reconnection. Although both in situ and ground-based observations provide evidence for steady and impulsive reconnection, the conditions governing when and where each occur remain unknown. Drake et al. (2006) suggest that antiparallel magnetosheath and magnetospheric magnetic fields favor steady reconnection along a single line, whereas shear angles less than \(127^{\circ}\) result in unsteady reconnection and the formation of magnetic islands or flux ropes. Steady reconnection predicts a gradual inward motion of the dayside magnetopause following southward IMF turnings, perhaps several Earth radii over a period of one to two hours (e.g., Aubry et al. 1970). It is not yet known how or whether the rate of this steady erosion changes with time. By contrast, sporadic reconnection models predict a sequence of abrupt earthward leaps, perhaps once each 8 minutes or so, corresponding to the equatorward jumps seen in ground-based radar and optical observations of the cusps when the IMF turns southward (Lockwood et al. 1989; Sandholt et al. 1998). Figure 5 shows one such sequence of events reported in ground-based observations of auroral emissions at 557.7 nm and 630.0 nm (Oksavik et al. 2005). With simultaneous solar wind observations, one can use soft X-ray observations to determine whether (Lockwood and Wild 1993) or not (Le et al. 1993) the bursts of reconnection corresponding to the inward magnetopause leaps are triggered by intrinsic magnetopause instabilities or fluctuations in the IMF orientation.

Fig. 5
figure 5

Meridian scanning photometer measurements from Svalbard (adapted from Oksavik et al. 2005). The top panel presents the 630.0 nm line while the bottom panel is the 557.7 line. Periodic poleward moving enhancements are observed

As a corollary, global images of the magnetopause location can be used to determine the time scale required for the magnetopause to move outward following a substorm onset or a northward IMF turning, and the mechanisms by which it does so. The outward motion of the magnetopause under these circumstances implies an addition of magnetic flux to the dayside magnetopause. The flux might be added by appending magnetosheath field lines to the dayside magnetosphere via either steady or unsteady simultaneous reconnection poleward of both cusps (Song and Russell 1992). Alternatively, the flux might be returned by sunward convection within the magnetosphere that continues even when the IMF turns northward (e.g., Øieroset et al. 1997). The rate of flux accretion remains unknown, but could be determined by tracking outward dayside magnetopause motion during intervals of northward IMF orientation.

Global perspectives can provide important information about the location and extent of reconnection. Component reconnection models predict erosion of the magnetopause along a tilted line passing through and centered on the subsolar point (Gonzalez and Mozer 1974; Laitinen et al. 2007). For many IMF orientations, antiparallel reconnection models predict reconnection at locations far from the subsolar point (Crooker 1979; Sandholt and Farrugia 2003). In the former case, magnetopause motion should begin at the subsolar point, in the latter case it should begin at locations away from the subsolar point. The manner in which reconnection spreads must also be determined. It could be initiated simultaneously over a wide region of the magnetopause as inferred from the sudden appearance of transient events in the high-latitude dayside auroral ionosphere (e.g., Lockwood et al. 1990), spread in the direction of the current at the speed of the current carriers for weak guide fields (Lapenta et al. 2006), or spread both along and opposite the current simultaneously at the Alfvén velocity for strong guide fields (Shepherd and Cassak 2012).

The ultimate extent of the reconnection line must also be determined. In some models, reconnection is very localized (Russell and Elphic 1979). In others, both steady and sporadic reconnection occur along reconnection lines that extend over many hours in local time (Lockwood et al. 1990; Phan et al. 2000). A small amount of localized plasmaspheric mass-loading may redistribute the locations where reconnection occurs on the magnetopause, whereas large mass loading might cause system level reconfigurations (Zhang et al. 2016). Finally, reconnection may also occur simultaneously at numerous sites spread across broad regions of the dayside magnetopause (e.g., Alexeev et al. 1998), in which case different portions of the magnetopause might erode inward erratically in a disjointed manner. Distinguishing between these possibilities requires global images of the magnetopause.

Inferences concerning the location and thickness of plasma boundary layers just inside the magnetopause can also provide information concerning the location of magnetopause reconnection. Wave-driven diffusion, reconnection facilitated by nonlinear Kelvin-Helmholtz instabilities, and reconnection at remote locations can all produce such boundary layers (e.g., Nakamura et al. 2006; Hasegawa et al. 2009). By contrast to diffusion, which generates boundary layers whose thickness increases with distance downstream from the subsolar point, or Kelvin-Helmholtz instabilities, which generate boundary layers whose thickness depends on downstream distances and magnetopause velocity shear (i.e., solar wind velocity), reconnection produces boundary layers of mixed magnetosheath and magnetospheric plasma whose width increases with distance from the reconnection site (Sonnerup et al. 1981; Gosling et al. 1990). In soft X-ray images the presence of these boundary layers will be detected as a blurring of the plasma boundaries that would otherwise be present. The presence and absence of these boundary layers can therefore be used to determine when and where reconnection is occurring, thereby distinguishing between component, antiparallel, and other reconnection models, each of which predicts a distinctly different reconnection location as a function of solar wind conditions.

We know very little about what influence other solar wind parameters such as the Mach number, plasma beta, or solar wind dynamic pressure have upon the rate and mode of reconnection, but this could be readily discerned from both detailed case and statistical studies of magnetopause erosion employing global observations of the magnetopause location and motion for different combinations of solar wind parameters. For example, there are reasons to suppose that reconnection, magnetic flux erosion, and the cross polar cap potential drop all saturate for strong southward IMF orientations (Mühlbachler et al. 2005). Global simulations indicate a slowdown and stall in dayside magnetopause erosion, overdraped lobes that extend further sunward than the dayside magnetopause, less magnetotail flaring than would be expected based on an extrapolation of empirical models, and even a dimple on the subsolar magnetopause for large negative IMF \(B_{z}\) (Dmitriev and Suvorova 2000, 2012; Siscoe et al. 2004; Ober et al. 2002, 2006), all features that should be readily seen in global images. Thus, global images could be used to determine the precise combination of solar wind parameters (e.g., dynamic pressure and IMF \(B_{z}\)) when saturation sets in (Yang et al. 2003).

Elsen and Winglee (1997) predicted that the location of the subsolar magnetopause would exhibit a diminished response to IMF \(B_{z}\) as the solar wind pressure increases, and a diminished response to solar wind dynamic pressure as the southward component of IMF \(B_{z}\) increases. Using the limited in situ observations available for unusual combinations of solar wind parameters, both case (Shue et al. 1998, 2001) and statistical (Roelof and Sibeck 1993; Lin et al. 2010) studies suggest that erosion is indeed non-linear, i.e. that the radial distance to the dayside magnetopause does not vary linearly with IMF \(B_{z}\), that the rate of erosion by IMF \(B_{z}\) diminishes for high solar wind dynamic pressures, and that the rate at which pressure changes compress the magnetosphere diminishes for strong southward IMF \(B_{z}\). Global images could confirm, extend, and quantify these results. Since the magnetopause does not respond instantaneously to variations in the IMF orientation (or the solar wind dynamic pressure) it will almost certainly be necessary to include the time history of the IMF orientation in determinations of magnetopause location (Shue et al. 2000).

Images can also be used to identify the degree to which radial IMF orientations reduce pressure upon the dayside magnetosphere (Fairfield et al. 1990) and allow the dayside magnetopause to expand outward (Merka et al. 2003b; Suvorova et al. 2010; Dušík et al. 2010), perhaps in response to kinetic effects within the foreshock or to magnetohydrodynamic anisotropies (Samsonov et al. 2012, 2013, 2017). They can be used to detect the effects, if any, of dawn/dusk or spiral/orthospiral IMF orientations on the size and shape of the steady-state magnetosphere. Finally, although the waves (Kaufmann and Konradi 1969; Samsonov et al. 2015) generated by most solar wind discontinuities may sweep along the bow shock and magnetopause too rapidly to be tracked, soft X-ray images could be used to track the response of both boundaries to very oblique discontinuities, i.e., those which traverse the dayside magnetosphere very slowly because their normals lie nearly transverse to the Sun-Earth line (e.g., Takeuchi et al. 2002).

The magnetospheric magnetic field perturbations associated with the Region 2 and ring current systems enhance magnetic field strengths in the outer dayside magnetosphere and might therefore be expected to push the magnetopause outward (Schield 1969). Numerical simulations suggest that the subsolar magnetopause moves outward some 0.6 to 0.8 Earth radii when the ring current intensifies (Samsonov et al. 2016). However, theory (Tsyganenko and Sibeck 1994) and some empirical models (Petrinec and Russell 1993) indicate that the dayside magnetopause moves outward only slightly during intervals when the ring current is enhanced. Observations suggest that the duskside magnetopause may (Wrenn et al. 1981; McComas et al. 1993; Dmitriev et al. 2004, 2005, 2011; Dmitriev and Suvorova 2012) or may not (McComas et al. 1994) lie further from Earth than the dawnside magnetopause in response to an enhanced partial ring current.

2.1.2 The Earth’s Cusps

Reconnection opens formerly closed magnetospheric magnetic field lines and allows solar wind mass, energy, and momentum to flow into the magnetosphere along bundles of open magnetic field lines that map from the magnetopause down to the high latitude dayside ionosphere (Heikkila and Winningham 1971). Plasma densities on these cusp magnetic field lines are slightly less than those in the magnetosheath, but far greater than those in the adjacent magnetosphere (Lavraud et al. 2004; Walsh et al. 2016a). Furthermore, the cusps extend deep into regions of the exosphere where neutral densities are very high. Consequently, the cusps must be bright soft X-ray emitters.

Because observations of the cusp are already available from both in situ (Escoubet et al. 1992; Pitout et al. 2006) and ground-based (Lockwood et al. 1989; Pinnock et al. 1993; Sandholt et al. 1998) observatories, one might ask why global images are needed. One answer is that it is difficult to extract complete comprehensive views of cusp behavior from the intermittent snapshots of in situ measurements along the paths followed by rapidly moving spacecraft. Another is that the spatially-limited optical views of the low-altitude cusp provided from a handful of stations in the northern and southern hemisphere tell us little about the cusp at mid- or high-altitudes. Global soft X-ray images will provide a broader view, one that connects our knowledge of magnetopause phenomena to the features seen on the ground. This section examines the wealth of information that can be learned about the solar wind-magnetosphere interaction from soft X-ray observations of the location, dimensions, motion, and structure of the Earth’s cusps.

First consider the location of the cusps in local time. Both component and antiparallel reconnection predict reconnection along the equatorial magnetopause during intervals of strongly southward IMF orientation. Component reconnection may continue on the subsolar magnetopause during intervals of strong dawnward or duskward IMF orientation (Gonzalez and Mozer 1974), but antiparallel reconnection moves away from the subsolar point to off-equatorial locations (Crooker 1979). Although cusps produced by component reconnection may remain in place near local noon when the IMF has a dawnward or duskward IMF orientation, the antiparallel reconnection model predicts that duskward (dawnward) IMF orientations move the northern cusp duskward (dawnward) but the southern cusp dawnward (duskward). During periods of strong dawnward or duskward IMF orientation, reconnection may occur at both high and low latitudes, forming double cusps (Wing et al. 2001; Berchem et al. 2016). Soft X-ray observations of cusp locations will determine whether component or antiparallel reconnection prevails as a function of solar wind conditions.

Now consider the latitude of the cusps. During periods of southward IMF, enhanced reconnection rates on the dayside equatorial magnetopause cause the cusps to move \(\sim10^{\circ}\) equatorward (Burch 1973; Carbary and Meng 1986; Wing et al. 2001). In the absence of simultaneous measurements in both hemispheres, we might suppose that the northern and southern hemisphere cusps move in unison to similar geomagnetic latitudes when the IMF turns southward. However, there is plenty of evidence indicating that their latitudes differ (Candidi and Meng 1988), for reasons that remain unclear. Global images with simultaneous solar wind coverage will afford an unprecedented opportunity to address this topic.

The response of the cusps to northward IMF orientations also remains to be fully established. Newell et al. (1989) reported observations indicating that reconnection moves to locations poleward of the cusp and appends magnetosheath magnetic field lines to the magnetosphere during periods of northward IMF orientation, causing the cusps to move poleward. By contrast, Palmroth et al. (2001) presented observations indicating equatorward cusp motion during intervals of strongly northward IMF and suggested that this might result from intensified reconnection on the equatorial magnetopause. Other work indicates that the latitudinal position of the high-altitude cusp does not move, but rather remains stationary for increasingly northward IMF orientations (Merka et al. 2002). Soft X-ray images will determine the latitudes of the cusps in both hemispheres as a function of time and discriminate between proposed models.

As cusp motion indicates the net rate at which magnetic flux is transferred from the dayside to the nightside magnetosphere (or vice-versa), determining the response time of the cusp to changing solar wind conditions and the velocity at which it moves is important to understand the state of the solar wind-magnetosphere interaction and time the development of storms and substorms. Yet the time scale for the cusp to respond to varying solar wind conditions remains unclear. Past observations indicate that the initial response begins almost immediately, but that a further 10 to 40 minutes are required to complete cusp relocations (Escoubet and Bosqued 1989; Němeček and Šafránková 2008). Yeoman et al. (2002) employed ground-based radar observations to track equatorward motion of the cusp during intervals of southward IMF orientation, but found no motion during intervals of northward IMF orientation. Pitout et al. (2006) reported several case studies in which snapshots from multipoint in situ observations indicated equatorward motion following southward IMF turnings, but poleward motion following northward IMF turnings. A wide field-of-view soft X-ray telescope will provide the sequences of images needed to identify cusp motion and time its velocity as a function of solar wind conditions. The observations could be used to determine whether or not steady-state conditions are ever achieved, how the magnetosphere responds to the onset of dayside and magnetotail reconnection, and how the magnetosphere responds to the cessation of dayside reconnection. Since the cusps lie at the boundary between open and closed magnetic field lines, observations of their latitude can immediately be used to quantify flux erosion from the dayside magnetosphere.

Just as in the case of the magnetopause, cusp motion can be steady, occur by leaps in response to individual southward IMF turnings (e.g., Lockwood et al. 1989), or occur by leaps in response to bursts of reconnection triggered by local magnetopause instabilities. The equatorward motion of the cusps may saturate for large southward IMF orientations (e.g., Siscoe et al. 2002; Ober et al. 2006). Little information is available concerning how the cusp moves in response to northward IMF turnings.

Now consider the response of the cusp to variations in the dipole tilt. Empirical models and both low- and high-altitude observations indicate that the cusps move equatorward in response to sunward diurnal and seasonal dipole tilts (Newell and Meng 1989; Zhou et al. 1999; Tsyganenko and Russell 1999). Both the width of the summer cusp and the densities within it exceed those of the winter cusp (Newell and Meng 1988; Pitout et al. 2006; Wiltberger et al. 2009). Simultaneous soft X-ray images of both cusps can be used to study these variations on a routine basis for the full range of solar wind and geomagnetic conditions, thereby quantifying how much plasma enters the magnetosphere in each hemisphere.

The width of the cusp yields important information concerning magnetospheric convection. The cusps span several Earth radii near the magnetopause (Walsh et al. 2012a) but narrow to dimensions of several hundred kilometers at their high-latitude, low-altitude, ionospheric footprints (Newell and Meng 1992). We adopt a kinetic interpretation to understand the internal structure of the cusps. The suprathermal magnetosheath particles entering the cusps precipitate into the high-latitude dayside ionosphere first, followed by the bulk of the distribution, and then the slower moving subthermal particles. Since the reconnected magnetic field lines within the cusps move in response to pressure gradient and magnetic field curvature forces, the precipitating particles exhibit distinctive spatial dispersion patterns (Rosenbauer et al. 1975; Reiff et al. 1977; Wing et al. 1996, 2001). The motion of magnetic field lines poleward from reconnection sites on the dayside equatorial magnetopause results in precipitating thermal and subthermal particle fluxes that initially increase abruptly and then subsequently decrease more gradually with latitude during periods of southward IMF orientation, as illustrated in Fig. 6. The width of the region over which they precipitate increases with increasing convection velocity. By contrast, during periods of northward IMF orientation, newly reconnected magnetic field lines either stagnate or move equatorward. Precipitating particle fluxes should either increase with latitude or show little variation. During periods of dawnward or duskward IMF orientation, curvature and pressure gradient forces should pull the newly reconnected magnetic field lines azimuthally, resulting in dawn/dusk cusp particle dispersion patterns. All these features, and their time-dependencies, could readily be identified and quantified by a global imager.

Fig. 6
figure 6

Cluster 4 CIS instrument measurements of density structure in the cusp. The spacecraft cuts through the high altitude cusp from low to high latitudes, i.e., from GSM \((R,\lambda) = (4.62~\mbox{R}_{E},54.5^{\circ})\) at 15:10 UT to (\(4.84~\mbox{R}_{E}, 64.5^{\circ}\)) at 1530 UT. The plasma density peaks at the equatorward edge and gradually decreases with increasing latitude. Here \(R\) is the radial distance from Earth and \(\lambda\) is the latitude

The azimuthal extent of the cusp in the direction transverse to the convection velocity provides information concerning the extent of the reconnection line(s) on the dayside magnetopause. Broad cusps may map to a line \(25~\mbox{R}_{E}\) long on the magnetopause for southward IMF orientations, but narrow cusps to a line only \(\sim5~\mbox{R}_{E}\) long for northward IMF orientations (Fuselier et al. 2002). Azimuthal structure within the cusp can be interpreted as evidence for patchy reconnection on the dayside magnetopause. If reconnection occurs simultaneously along a single extended reconnection line, cusp properties will vary smoothly in azimuth. Whether or not it occurs simultaneously, patchy reconnection along multiple disconnected reconnection line segments will result in considerable azimuthal structure. Images of the cusp will provide information concerning the extent of reconnection on the dayside magnetopause.

Steady reconnection along a single reconnection line for either southward or northward IMF orientations should produce smooth variations in ion energy and density versus latitude. Stepped structures in the meridional direction (Newell and Meng 1991; Escoubet et al. 1992; Trattner et al. 2008) can therefore be interpreted as evidence either for time-varying reconnection (Smith and Lockwood 1990; Escoubet et al. 1992) or multiple reconnection sites at different latitudes (Kan 1988; Nishida 1989; Onsager et al. 1995; Trattner et al. 1999). Spatial and temporal variations can occur at the same time (Němeček et al. 2004). Soft X-ray images can be used to distinguish between these possibilities. Steady-state structures generated by multiple reconnection sites remain in place, whereas transient features produced by time-dependent reconnection convect antisunward. Images could also be used to determine the number and extent of such features, thereby addressing the locations of reconnection and the relative importance of steady and transient reconnection.

Finally, just as in the case of the magnetosphere as a whole, an increase in the solar wind dynamic pressure may diminish the dimensions of the cusp (Fung 1997). However, studies indicate that an increase in the solar wind dynamic pressure causes the dimensions of the cusp to increase (Zhou et al. 2000; Merka et al. 2002). Simulation results suggest that cusp dimensions initially increase with increasing solar wind pressure, but saturate near solar wind dynamic pressures of 3 nPa (Zhang et al. 2013). Perhaps the cusp widening results from greater magnetosheath magnetic field strengths and reconnection rates during intervals of enhanced solar wind dynamic pressure magnetopause (Newell and Meng 1994).

2.1.3 Transients at Earth’s Magnetopause and in the Cusps

Transient structures/events with durations on the order of 30 s to several minutes are common in the vicinity of the Earth’s magnetopause. They have been interpreted as the magnetospheric response to variations in the intrinsic solar wind dynamic pressure (Kaufmann and Konradi 1969), the magnetospheric response to transient dynamic pressure fluctuations generated within the foreshock (Fairfield et al. 1990), the Kelvin-Helmholtz instability operating at the magnetopause (Boller and Stolov 1973), and flux transfer events (FTEs) generated by bursts of magnetic reconnection between magnetosheath and magnetospheric magnetic field lines (Russell and Elphic 1978). If sufficiently numerous and extensive, the events might contribute significantly to (Lockwood et al. 1990) or even dominate (Lockwood et al. 1995) the solar wind-magnetosphere interaction. Consequently, quantifying the significance of each proposed transient solar wind-magnetosphere interaction mechanism as a function of solar wind conditions is a core objective of magnetospheric physics.

Although comprehensive single point and multipoint in situ measurements provide evidence for each of the proposed mechanisms, only instantaneous global measurements can definitively quantify their significance on the basis of their occurrence rates and dimensions. Fortunately, models for the various transient interaction mechanisms make very specific predictions concerning event occurrence patterns and signatures.

Solar wind tangential discontinuities are relatively common, arriving at Earth about once per hour (Burlaga and Ness 1969). Very few tangential discontinuities provide density variations greater than 35% (Solodyna et al. 1977). Although much rarer, interplanetary shocks often provide factor of two or larger density and dynamic pressure variations (e.g., Wang et al. 2010). Because they extend over many Earth radii transverse to the Sun-Earth line (Burlaga and Ness 1969), the pressure variations that accompany solar wind discontinuities launch widespread antisunward moving waves on the magnetopause. Transient enhancements in the solar wind dynamic pressure compress the magnetopause, while transient decreases allow it to expand outward. The same discontinuities launch fast mode waves that propagate throughout the magnetosphere. These fast mode waves may outrun the antisunward-moving solar wind discontinuities and initiate magnetopause motion ahead of the driving solar wind discontinuities. For example, the fast mode compressional waves launched by a transient increase in the solar wind dynamic pressure may cause the magnetopause to move outward in advance of the inward motion associated with the discontinuity itself (Kaufmann and Konradi 1969; Samsonov et al. 2015). The extent and amplitude of pressure-pulse induced waves could be determined by correlating global images of magnetopause motion with simultaneous in situ observations of solar wind dynamic pressure.

Kinetic effects in the foreshock generate more localized density and pressure variations (Thomas and Brecht 1988; Omidi and Sibeck 2007). Some (e.g. hot flow anomalies) lie centered on tangential discontinuities, others (e.g., foreshock cavities) are bounded by tangential discontinuities (Sibeck et al. 2001), some (e.g., compressional boundaries) bound the foreshock (Omidi et al. 2009), and yet others (e.g., spontaneous hot flow anomalies) lie within the foreshock but are not associated with discontinuities (Zhang et al. 2013). Corresponding ripples in the bow shock position result in magnetosheath plasma jets with enhanced densities capable of driving transient magnetopause motion and magnetospheric compressions (Hietala et al. 2012). The impact of these events on the magnetosphere should be greatest during intervals of radial or near-radial IMF orientation, as illustrated in the left panel of Fig. 7, when the IMF lies nearly along the Sun-Earth line and the foreshock lies upstream from the Earth’s dayside magnetosphere (Fairfield et al. 1990). Because the foreshock lies upstream from the pre-noon bow shock and dayside magnetosphere for the typical spiral IMF orientation (see the right panel of Fig. 7), the magnetopause boundary waves and fast move waves transmitted into the magnetosphere by foreshock pressure pulses should generally be limited to the pre-noon magnetosphere (e.g., Howe and Binsack 1972; Rufenach et al. 1989; Russell et al. 1997). In situ observations indicate that the foreshock pressure pulses are more prominent during intervals of enhanced solar wind velocities (Sibeck et al. 2001; Facskó et al. 2008). Consequently, we expect the same to be true for the corresponding magnetopause motion. The significance of foreshock events can be determined by combining global images of magnetopause motion with in situ observations of solar wind variations, and in particular IMF orientations.

Fig. 7
figure 7

Density structures resulting from kinetic processes within the foreshock (adapted from von Alfthan et al. 2014). The panels display density (\(\mbox{cm}^{-3}\)) from Vlasiator code hybrid-Vlasov simulations. Solar wind parameters are identical for the two panels, with the exception of the IMF orientation, which is radial in panel (a) but inclined \(30^{\circ}\) from radial in panel (b). The white line is parallel to the IMF orientation

A Kelvin-Helmholtz instability occurs when flow shears at the magnetopause or inner edge of the low-latitude boundary layer overcome stabilizing curvature forces in draped magnetosheath and magnetospheric magnetic field lines and generate antisunward-propagating/convecting waves. The fastest growing wavelengths should be about 10 times greater than boundary layer thicknesses, with wave amplitudes increasing with increasing shears (Walker 1981) and downstream distance (Li et al. 2012). The instability is most likely to occur when strong flow shears lie perpendicular to both magnetosheath and magnetospheric magnetic field orientations, a condition most readily obtained on the equatorial flanks of the magnetosphere during intervals of strongly northward or southward IMF orientations (Southwood 1968). However, the instability can occur at other locations, including the high latitude magnetopause, when conditions are favorable (Hwang et al. 2012). It may be very common. Kelvin-Helmholtz waves occur about 40% of the time when the IMF points northward and about 10% of the time when it points southward. (Kavosi and Raeder 2015). Weaker magnetosheath magnetic field components parallel to the flow shear may make the instability more likely on the side of the magnetosphere behind the quasi-parallel bow shock (Nykyri 2013). We do not know if conditions sometimes favor hemispheric asymmetries in the occurrence of Kelvin-Helmholtz waves (Taylor et al. 2012) These predictions of the Kelvin-Helmholtz model can be tested by comparing global observations of magnetopause motion with in situ measurements of solar wind parameters.

Flux transfer events (FTEs) are bundles of intertwined magnetic field lines that, in contrast to the boundary waves generated by pressure pulses and the Kelvin-Helmholtz instability, simultaneously bulge outward into both the magnetosheath and magnetosphere. They contain a mixture of magnetosheath and magnetospheric plasmas, and consequently locally broaden and diminish the otherwise sharp density gradients that mark the magnetopause. Erkaev et al. (2003) attribute abrupt, pronounced, decreases and gradual increases of the density in the inner magnetosheath to bursts of reconnection. Because they result from reconnection, and reconnection is more likely when and where the shear between magnetosheath and magnetospheric magnetic field orientations is greater, FTEs on the dayside magnetopause are more common during intervals of southward IMF orientation (Berchem and Russell 1984). The origin of events on the flanks of the magnetosphere remains disputed. They may be generated by tilted reconnection lines that extend antisunward from the subsolar magnetopause (Kawano and Russell 1997), or be generated locally in regions of the high-latitude magnetopause where magnetosheath and magnetospheric magnetic field lines lie antiparallel (Sibeck et al. 2005). The fate of FTEs remains equally uncertain. They may slip over the polar magnetopause or be destroyed by interactions with magnetospheric magnetic field lines within the cusp regions (Omidi and Sibeck 2007). The occurrence of transients and FTEs may, or may not, be triggered by the arrival of solar wind discontinuities (Le et al. 1993; Lockwood and Wild 1993; Tkachenko et al. 2011). These and other questions could be readily answered with simultaneous global images of the magnetopause and in situ solar wind observations.

2.1.4 Earth’s Magnetosheath

The magnetosheath envelops the magnetosphere, thereby providing its outer boundary conditions and the medium through which solar wind features are modified and transmitted to the magnetopause. Magnetosheath properties govern the occurrence patterns for reconnection and the Kelvin-Helmholtz instability at the magnetopause, which in turn control the flow of solar wind mass, energy, and momentum into the magnetosphere. In particular, low densities and low plasma beta favor the occurrence of magnetic reconnection (Phan et al. 2013), perhaps enabling steady reconnection to occur on high-latitude regions of the magnetopause where it would otherwise be precluded by high magnetosheath velocities (Fuselier et al. 2000; Avanov et al. 2001; Panov et al. 2008). By contrast, high densities and low Alfvén velocities favor the occurrence of the Kelvin-Helmholtz instability (e.g., Southwood 1968; Walsh et al. 2015).

Spreiter et al. (1966) reported the predictions of a gasdynamic model for an axially symmetric magnetosphere. Densities decrease slightly from the subsolar magnetopause to the bow shock along radial lines within \(\sim45^{\circ}\) from the Sun-Earth line, but increase significantly from the magnetopause to the bow shock along radial lines at greater angles. MHD theory suggests that the presence of a magnetic field within the flowing plasma results in the formation of a plasma depletion layer (PDL) with very low magnetosheath densities but enhanced magnetic field strengths just outside the dayside magnetopause (Zwan and Wolf 1976). Numerical simulations indicate that stable depletion layers are present during intervals of steady northward IMF orientation. It can be difficult for individual spacecraft to detect the predicted smooth transitions and non-uniform increases in layer thickness with both latitude and longitude away from the subsolar point due to the back and forth motion of the layers in response to constantly varying solar wind plasma parameters (Wang et al. 2003). On the other hand, X-ray imagers should readily identify the appearance and disappearance of a PDL as a change in emission intensity and width of the magnetosheath to magnetosphere transition.

Modeling case studies suggest that the PDL can extend to cusp latitudes and 6 hours in local time away from noon (Wang et al. 2003). Predicted density depletion factors (for similar solar wind conditions) range from 1.2 (Lyon 1994) to 10 (Siscoe et al. 2002). Table 1 summarizes reported depletion layer dependencies on solar wind conditions. In some models, depletion factors and layer thicknesses diminish with increases in the IMF clock angle away from northward in the plane perpendicular to the Sun-Earth line (Siscoe et al. 2002), while in others they increase (Table 5 of Wang et al. 2004a; Pudovkin et al. 1995, 2001, 2002). Wang et al. (2004a) presented results from a parametric study indicating that depletion layer widths decrease with increasing solar wind magnetosonic Mach number, increase with increasingly northward IMF Bz strength, increase as the (clock angle) component of the IMF in the plane perpendicular to the Sun-Earth line rotates away from due northward, and remain almost constant as the dipole tilt increases. Wang et al. (2004a) also concluded that density depletion factors increase and then decrease as the solar wind magnetosonic Mach number increases, increase but then decrease as with increasingly northward IMF \(B_{z}\) strength, increase slightly as the IMF clock angle increases, and decrease as the dipole tilt increases. Simulation results presented by Maynard et al. (2004) indicate that depletion layers form just outside the dayside magnetopause even for southward IMF orientations when the IMF has a finite component along the Sun-Earth line and/or there is a strong dipole tilt, because reconnection moves to higher latitudes. Furthermore, they indicate depletion layers forming poleward of the cusps during intervals of very strongly southward IMF orientations.

Table 1 Depletion layer predictions and observations

Case and statistical studies of single-point in situ spacecraft observations provide support for some of these predictions. On the subsolar magnetopause, both layer thicknesses and depletion factors diminish when the IMF turns southward (Phan et al. 1994; Slivka et al. 2015) or radial (Anderson and Fuselier 1993). Farrugia et al. (1995) reported that the layer becomes more pronounced for low solar wind Mach numbers, while Anderson et al. (1997) reported a pronounced layer for high solar wind Mach numbers when the magnetosphere is compressed by high solar wind dynamic pressures, solar wind densities are large, and Alfvén velocities are low. Maynard et al. (2004) reported that the layer shifts to the region behind the quasi-perpendicular bow shock. Finally, a pronounced depletion layer has indeed been observed on the high latitude magnetopause during an interval of southward IMF orientation (Moretto et al. 2005). Contrary to model predictions, the depletion layer may become less prominent for small IMF cone angles (Anderson and Fuselier 1993). Soft X-ray images, like those proposed in this work, could be used to discriminate between these predictions and examine others yet to be tested.

Song et al. (1990) and Song and Russell (1992) reported observations of anticorrelated density enhancements and magnetic field strength depressions just upstream from the subsolar magnetopause and interpreted these observations in terms of standing slow mode waves. Southwood and Kivelson (1992, 1995) illustrated how a slow mode wave standing in the magnetosheath could result in a region with enhanced densities and depressed magnetic field strengths. Magnetic field lines within this region would have greater components parallel to the Sun-Earth line than those either further upstream in the magnetosheath proper or downstream in the depletion layer. Lee et al. (1991) identified the anticorrelated features in two-dimensional incompressible MHD simulations whenever there was a magnetic field component parallel to the Sun-Earth line. However, Wang et al. (2004b,c) and Samsonov and Hubert (2004) were unable to find any such features in global MHD simulations for any IMF orientation. Hubert and Samsonov (2004) concluded that the anticorrelated density enhancements and magnetic field strength decreases were simply antisunward propagating solar wind features caught just before they encountered the magnetopause, which prompted a comment (Song et al. 2005) and reply (Hubert and Samsonov 2005). The issue remains unsettled, but could be addressed by imaging the structure of the inner magnetosheath.

Dawn/dusk asymmetries in magnetosheath densities may control the occurrence of reconnection and the entry of solar wind/magnetosheath plasma into the magnetosphere. This entry results in the formation of low-latitude boundary layers with magnetosheath-like plasma at densities lower than those in the magnetosheath. Observations indicating greater densities in the dawnside than duskside low-latitude boundary layer (LLBL, Hasegawa et al. 2003) and magnetotail plasma sheet (Wing et al. 2005) suggest greater pre- than post-noon magnetosheath densities and/or plasma entry. Walters (1964) argued that the presence of a Parker spiral IMF embedded in the flowing solar wind plasma would indeed result in greater dawnside than duskside magnetosheath densities, particularly during intervals of low solar wind Mach number. Global MHD models confirm this prediction for spiral IMF orientations, with asymmetries that increase for decreasing solar wind Mach number (Walsh et al. 2012b). Observationally, Paularena et al. (2001), Němeček et al. (2002), and Longmore et al. (2005) report greater densities in the dawnside magnetosheath than in the duskside magnetosheath. However, each of these studies concluded that the density asymmetry was unrelated to the IMF orientation. By contrast, Walsh et al. (2012b) reported asymmetries in the expected sense. A statistical survey reported by Dimmock and Nykyri (2013) found no evidence for any dawn/dusk density asymmetry, but Dimmock et al. (2016) went on to show that the expected asymmetries were in fact present, but only in the region immediately outside the magnetopause. Finally, note that greater densities and consequently enhanced plasma betas should inhibit reconnection. Rather than resulting from asymmetric magnetosheath densities, observations of enhanced densities in the dawnside boundary layer and plasma sheet may indicate the preferential operation of one or more diffusive entry processes.

2.1.5 Earth’s Bow Shock

Simulations for the solar wind’s interaction with the Earth’s magnetosphere require accurate values for the polytropic index \(\gamma\) which represents the ratio of specific heats (\(C_{p}/C_{v}\)) and closes the set of magnetohydrodynamic equations. Determining the polytropic index is important because it controls phenomena as diverse as the degree of heating in magnetic reconnection (Hesse and Birn 1992) and magnetosheath flow deflections (Nishino et al. 2008). Theoretical values for \(\gamma\) range from 2 (for an adiabatic gas with two degrees of freedom perpendicular to the magnetic field), through 5/3 (for an adiabatic gas with three degrees of freedom), 1.5, and 1.33 (when there is a heat flux escaping from the magnetosheath into the solar wind, Nishino et al. 2008), to 1 (for an isothermal gas). Observationally-inferred values for \(\gamma\) are almost as diverse, ranging from 1.67 (Russell et al. 1983), through 1.76 (Farris et al. 1991) and 1.85 (Tatrallyay et al. 1984), to 2 (Zhuang and Russell 1981).

Density jumps at the bow shock provide crucial information concerning \(\gamma\) (Farris et al. 1991). Following Spreiter et al. (1966), the jumps are a function of both \(\gamma\) and the upstream solar wind magnetosonic Mach number (MMS), i.e., \(\rho/\rho_{{sw}}=(\gamma+1){\mathrm{M}}^{2}_{\mathrm{MS}}/ [(\gamma- 1){\mathrm{M}}^{2}_{\mathrm{MS}}+ 2]\), where \(\rho\) is the density in the subsolar magnetosheath. For typical values of \(\mathrm{M}_{\mathrm{MS}}\gg1\), \(\rho/\rho_{{sw}}\) approaches \(\mathrm{M}^{2}_{\mathrm{MS}}\) for \(\gamma=1\), 4 for \(\gamma={5\over 3}\), and 3 for \(\gamma=2\). Alternatively, the locations of the bow shock and magnetopause themselves can also be used to determine \(\gamma\). As noted by Farris et al. (1991), \(\gamma=[(1.1+\Delta/D){\mathrm{M}}^{2}_{\mathrm{MS}}-2.2]/[(1.1-\Delta/D){\mathrm{M}}^{2}_{\mathrm{MS}}]\), where \(D\) is the standoff distance of the magnetopause from the center of the Earth and \(\Delta\) is the distance between the bow shock and the magnetopause. There are alternative formulations within the gasdynamic framework, including those that take into account the possibility that the density jump approaches unity and the bow shock recedes to infinity as the Mach number approaches unity, or that the solar wind feels the effects of the magnetospheric shape rather than the distance between the bow shock and the magnetosheath (Farris and Russell 1994). There are also magnetohydrodynamic approaches (Cairns and Grabbe 1994; Grabbe and Cairns 1995). When the Mach number approaches unity, Alfvén wings may form, greatly modifying the size and shape of the magnetopause (Ridley 2007; Chané et al. 2012).

The various models make strikingly different predictions for the location of the subsolar bow shock as a function of IMF orientation and solar wind Mach number. Cairns and Lyon (1996) predicted that the standoff distance increases as the solar wind Mach number decreases for IMF orientations transverse to the Sun-Earth line, but decreases for IMF orientations parallel to the Sun-Earth line. Models presented by Cairns and Grabbe (1994) and Cairns and Lyon (1996) predict standoff distances for low solar wind Mach numbers far greater than those predicted by Verigin et al. (2001) or Farris and Russell (1994). As illustrated in Fig. 8, the quasi-perpendicular bow shock lies further upstream than the quasi-parallel bow shock, with the discrepancy increasing as the solar wind Mach number decreases (Chapman and Cairns 2004). The latter authors predict a dimple on the subsolar bow shock for very low solar wind Mach numbers and radial IMF orientations.

Fig. 8
figure 8

Bow shock locations and shapes predicted by MHD models for three solar wind Alfvénic Mach numbers (adapted from Chapman et al. 2004). The IMF lies \(45^{\circ}\) from the Earth-Sun line

It has proven difficult to verify these predictions with studies employing in situ observations. Despite multipoint observations, Fairfield et al. (2001) was unable to discriminate between the models for an unusually distant bow shock for low solar wind Mach numbers. Consistent with expectations, Slavin et al. (1996), Merka et al. (2003b), and Jelínek et al. (2010) reported subsolar bow shock locations closer to Earth and therefore very thin subsolar magnetosheaths during intervals of radial IMF orientation. Verigin et al. (2001) reported results from a small statistical study indicating that the standoff distance to the bow shock increases with increasing Alfvénic Mach number for field-aligned solar wind flows, but decreases for non-field-aligned flows. However, Jeřáb et al. (2005) could find no dependence of the bow shock location upon the IMF orientation whatsoever. Jeřáb et al. (2005) attributed the absence of any inward bow shock motion associated with southward IMF turnings and inward magnetopause erosion to a compensatory increase in the magnetosheath thickness associated with the blunter obstacle posed by an eroded dayside magnetopause. Nevertheless, Jeřáb et al. (2005) did find that the distance to the bow shock increases linearly as a function of the IMF strength.

One might expect abrupt variations in solar wind parameters to drive corresponding inward and outward motions of the bow shock and magnetopause (e.g., Fairfield et al. 2001). Consistent with this hypothesis, Anderson et al. (1968) found a good correspondence between the periods and amplitudes of bow shock and magnetopause motion on individual spacecraft passes. However, the similarity of the periods and amplitudes does not necessarily mean the two boundaries move inward and outward in phase. Korotova et al. (2012) recently reported that the same change in the IMF orientation drove transient outward bow shock motion, but transient inward magnetopause motion. And Jelínek et al. (2006) reported that the motion of the bow shock does not correspond to that of the magnetopause.

Summarizing results to date, Merka et al. (2003a, 2005) noted that existing models for the bow shock underestimate the distance to the bow shock under strong IMF conditions, fail to reflect the effects of variations in the IMF and solar wind velocity vectors, and do not correctly describe the bow shock location during intervals of low solar wind Mach number. Even large statistical studies based on in situ observations fail to resolve expected dawn/dusk differences and Mach cone asymmetries. Global images of the bow shock and magnetopause should be able to resolve these and other issues by identifying the locations of the bow shock and magnetopause, determining the density jump at the bow shock, discriminating between models, and providing the information needed to determine \(\gamma\).

2.1.6 Earth’s Foreshock

The magnitude of the jump in magnetic field strengths (or densities) at the bow shock determines its ability to accelerate particles. Shock-drift acceleration at the quasi-perpendicular bow shock produces beams of ions and electrons on magnetic field lines that lie perpendicular to the bow shock normal (Decker 1983). The maximum energy gained by the reflected particles is given by \({T_{f}/T_{i}}=2r[1+(1-r^{-1})^{1\over 2}]-1\), where \({T_{i}}\) and \({T_{f}}\) are the initial and final particle energies and \(r\) is the ratio of the magnetosheath to IMF strengths. Solar wind ions with \(\sim1~\mbox{keV}\) energies might be accelerated to \(\sim14~\mbox{keV}\) for \(r=4\). By contrast Fermi acceleration of an incident monoenergetic particle distribution at the quasi-parallel bow shock can produce diffuse ion populations with far greater energies, near-isotropic pitch angle distributions, and power law spectra whose spectral indices depend upon the ratio of magnetosheath to interplanetary magnetic field strengths. In the non-relativistic case, the steady-state spectral index for the distribution function is given by \(3r/(r-1)\) (Blandford and Ostriker 1978).

Soft X-ray images can provide the information needed to determine the extent and nature of particle acceleration at the bow shock. First, the images can be used to identify the transition between the quasi-parallel and quasi-perpendicular bow shocks, which is expected to occur where the angle between the IMF and the normal to the bow shock passes through \(45^{\circ}\). A sharp density discontinuity indicates the quasi-perpendicular bow shock, whereas a broader and far more turbulent transition should mark the quasi-parallel bow shock. Secondly, the images can be used to determine the strength of the density, and consequently the magnetic field strength, jump at the bow shock.

Kinetic effects generate a wealth of mesoscale density structures upstream from Earth’s bow shock, including hot flow anomalies (Thomsen et al. 1986), foreshock cavities (Sibeck et al. 2001), density holes (Parks et al. 2006), and bubbles (Turner et al. 2013). By enhancing and/or diminishing upstream densities, deflecting solar wind flows, and perturbing corresponding magnetosheath parameters, these structures generate prominent transient events in the outer dayside magnetosphere and dayside auroral ionosphere. However, with one exception, the limited dimensions and ephemeral nature of most of these features and their magnetospheric responses probably preclude soft X-ray imaging. The exception is the foreshock compressional boundary, a region of enhanced density piled up on the edges of the quasi-parallel foreshock (Omidi et al. 2009, 2013). Numerical simulations indicate that these structures can be quasi-steady-state features for a wide variety of IMF orientations. Some observational studies support this point of view, while others interpret the density enhancements as foreshock-generated structures moving antisunward with the solar wind flow (Sibeck et al. 2008; Billingham et al. 2008, 2011). Since the density enhancements and depletions associated with the structures extend nearly normal to the bow shock, it should be relatively easy to employ global images to distinguish between these two models.

2.2 Comets

Many visible light observations and a handful of in situ measurements provide tantalizing views of the complex plasma phenomena that occur when the solar wind encounters comets. As shown in Fig. 9, in situ observations indicate that these structures include a bow shock, a “cometopause”, and an ionopause (Mendis 1988; Flammer 1991; Mendis and Horányi 2014). The bow shock forms in response to mass loading. As they approach the Sun, comets sublimate large clouds of neutral gas. The solar wind flow picks up ionized atoms and molecules within this cloud. If sufficiently numerous, the pick-up ions slow the flow down to the point where a bow shock forms. Deeper inside the bow shock, a “collisionopause” or “cometopause” forms at the transition from the heated and decelerated shocked collisionless mass-loaded flow to flow cooled and even more significantly decelerated by collisions and charge exchange with expanding cometary neutral molecules, in addition to pick-up ions generated by photoionization. Still closer to the nucleus lies the ionopause, the locus of points where the solar wind plasma makes its closest approach to the comet.

Fig. 9
figure 9

Schematic representation of the global morphology of the solar wind interaction with a cometary atmosphere, showing the various discontinuities in the flow pattern (adapted from Mendis 1988)

In situ observations confirm the presence of weak shocks on the nightside flanks of cometary tails (Coates 1995). Dayside shock strengths, and corresponding density enhancements, should be far greater. Simulations demonstrate that the IMF orientation controls the nature and thickness of the bow shock (Omidi and Winske 1987) which is thin for quasi-perpendicular configurations, broader for an intermediate shock, and narrow again for quasi-parallel shocks. Theory indicates that the dimensions of the bow shock increase as the solar wind Mach number diminishes and/or gas production rate increases. Recent models show that the dimensions of the bowshock increase with greater photoionization, charge-exchange, and electron impact ionization (Simon Wedlund et al. 2017). Consequently, the dimensions of the magnetosheath region of shocked solar wind plasma behind the dayside bow shock should also increase as comets move sunward and sublimation increases (Flammer 1991). Furthermore, since the rate of ion pick up via microscopic wave-particle interactions should be lower than that via macroscopic motional \(\mathbf{E}\times\mathbf{B}\) electric fields, a quasi-parallel bow shock should lie nearer the comet than a quasi-perpendicular bow shock (Omidi and Winske 1986).

The cometopause separates fast moving shocked solar wind flow from a region dominated by compressed IMFs and cometary ions. The width of the cometopause, where solar wind densities diminish, may be abrupt (\(\sim10^{4}~\mbox{km}\)) (Gringauz et al. 1986) perhaps in response to a charge exchange avalanche (Gombosi 1987), or more diffuse (Balsiger et al. 1986; Amata et al. 1986), with the width depending upon the IMF orientation (Galeev et al. 1988). Within the cometopause, the solar wind proton flow decelerates rapidly and cools in response to charge exchange with cometary neutrals. Correspondingly, the densities of both major (\(\mbox{H}^{+}\), \(\mbox{He}^{++}\)) and minor (e.g., \(\mbox{O}^{6+}\), \(\mbox{C}^{5+}\)) solar wind species should increase. Finally, no solar wind ions reach locations closer to the comet than the ionopause, although comets with sufficiently low outgassing rates may lack an ionopause. As in the case of the bow shock, theory predicts that the cometopause and ionopause structures will move outward as comets approach the Sun and sublimation increases (Flammer 1991).

Optical emission in cometary tails is dominated by molecular band emission due to \(\mbox{CO}^{+}\) and \(\mbox{H}_{2}\mbox{O}^{+}\)  while the in the cometary head, it is dominated by molecular band emission from \(\mbox{C}_{2}\) and the reflection of sunlight by dust. In contrast to optical observations, soft X-ray observations (e.g., Lisse et al. 1996; Dennerl et al. 1997; Gao and Kwong 2002) are dominated by interactions with water and its dissociation products OH, O, and H (Bodewits et al. 2007) and can be used to determine the characteristics of the dayside plasma and neutral environments of comets on a routine basis as a function of solar wind conditions and distance from the Sun. As illustrated in Fig. 10, theory predicts and observations confirm that the attenuation of solar wind ion densities via charge transfer collisions with increasing depth into the extended cometary atmosphere or coma results in integrated LOS soft X-ray emissions that peak in a bowl-shaped region within the magnetosheath on the sunward side of cometary nuclei (Wegmann et al. 2004). (However, there are cases such as 2P/Encke (Lisse et al. 2005), where the coma is of sufficiently low density that it is collisionally thin to charge-exchange, in which case the morphology is roughly spherical.) Soft X-ray emissions should be far greater in the magnetosheath than in the solar wind thanks to greatly enhanced plasma densities and thermal velocities in the magnetosheath, as well as greater neutral densities. Beyond the bow shock, emissions should fall off as an inverse function of radial distance from the nucleus. Both the intensity and dimensions of the emitting region depend on the rate of neutral gas production, and should therefore increase as comets approach the Sun. Individual line intensities also depend upon the flux of high charge state ions, which varies with the state of the solar wind and the concomitant ion abundances (Bodewits et al. 2007). Finally, the Kelvin-Helmholtz instability may locally permit solar wind plasma to penetrate deep into cometary ionospheres as predicted by Ershkovich and Mendis (1983) and as seemingly observed at 67P/Churyumov-Gerasimenko by Goetz et al. (2016a,b).

Fig. 10
figure 10

Shading shows the soft X-ray intensities of Comet Hyakutake observed by the ROSAT Wide Field Camera (WFC), while contours show the intensity of the best adapted hydrodynamic model (Wegmann et al. 2004). The nucleus lies at \((X,R)=(0,0)\), where \(X\) points towards the Sun. Lengths are in units of \(10^{5}~\mbox{km}\)

2.3 Mars and Venus

Identifying and quantifying the processes that cause atmospheric loss is a major objective of planetary studies. They were the principle objectives of NASA’s recent MAVEN mission to Mars (Jakosky et al. 2015b), and important objectives of ESA’s Mars (Chicarro et al. 2004) and Venus (Titov et al. 2006) Express missions.

A number of processes govern the loss of the Martian and Venusian atmospheres (e.g., Nagy et al. 2004; Lammer et al. 2006; Dubinin et al. 1996; Lundin 2011). Some invoke bombardment and hydrodynamic outflow. Others involve solar wind interactions with the planetary atmosphere and/or ionosphere, such as the removal of pick up ions generated by photoionization or charge exchange, or the formation of detached blobs of ionospheric plasma generated by either the Kelvin-Helmholtz instability at the ionopause or magnetic reconnection with ionospheric or remnant crustal magnetic fields. Even in the absence of solar wind stripping, ambipolar electric fields may cause a planet to lose heavy ions (Collinson et al. 2015, 2016).

Although there is evidence for enhanced escape of ionospheric ions during space weather storms (Luhmann et al. 2007; Edberg et al. 2011; Collinson et al. 2015; Jakosky et al. 2015a), assessing the significance of these and other mechanisms with isolated in situ measurements can be difficult. Models can help (Lillis et al. 2015), but models for Mars predict escape fluxes that differ by more than an order of magnitude (Brecht and Ledvina 2006; Brain et al. 2010b). Observations indicate similar variations over the course of a solar cycle (Lundin et al. 2013), which must be due at least in part to the large variations in exospheric densities that occur over the solar cycle (e.g., Forbes et al. 2008). The significance of each mechanism that invokes solar wind-planetary interactions can be quantified using global images of the corresponding diagnostic plasma density structures.

As in the case of comets, the interaction of the supersonic solar wind with Mars generates several plasma structures where the densities of ions with solar wind origin change abruptly (Brain 2006; Bößwetter et al. 2007). Some of these boundaries are illustrated in Fig. 11, while densities from a numerical simulation are shown in Fig. 1. The outer edge of the magnetosheath is bounded by a bow shock where densities increase abruptly from solar wind to magnetosheath values. The density of ions diminishes gradually from the magnetosheath to the ionospheric side of the magnetic pile-up region (MPR) on the inner edge of the magnetosheath, in a manner akin to that in the depletion layer outside Earth’s magnetopause. Electron observations indicate that this boundary is either lumpy or permeable in regions of radial crustal magnetic fields (Brain et al. 2005). Finally, the foreshock lies upstream from the bow shock on IMF lines connected to that boundary. By analogy to Earth, we expect regions of enhanced solar wind densities to bound a foreshock exhibiting depressed densities.

Fig. 11
figure 11

Cartoon of the global Martian solar wind interaction (Brain 2006). Orange shading indicates the density of planetary neutrals. Blue indicates the relative density of solar wind ions in different plasma regions (labeled in black), separated by different plasma boundaries (labeled in magenta). Here MPR stands for the magnetic pileup region, MPB for magnetic pile up boundary, and PEB for the photoelectron boundary

Table 2 (Brain 2006) summarizes the reported effects of the solar wind dynamic pressure, IMF orientation, solar extreme ultraviolet (EUV) radiation, season, and crustal magnetic fields on the distances to the bow shock, magnetic pile-up boundary, and photoelectron boundary at Mars. Increases in the solar wind dynamic pressure may (Brain et al. 2005; Crider et al. 2005; Morgan et al. 2014) or may not (Trotignon et al. 1996) move the magnetic pile-up and photoelectron boundaries towards the planet. The terminator bow shock lies further from the planet in the directions perpendicular to the IMF orientation (Zhang et al. 1991a). The Martian bow shock flares and moves further from the planet as the solar wind Mach number decreases (Edberg et al. 2010).

Table 2 Drivers affecting the variability of plasma boundary locations at Mars (Bertucci et al. 2005; Zhang et al. 1991a,b)

Crustal magnetic fields affect plasma and magnetic field structures in the vicinity of Mars. They raise the distances to the magnetic pile-up and photoelectron boundaries, thereby locally precluding direct solar wind interactions with the ionosphere, but they do not appear to increase the distance to the bow shock (Brain 2006, and references therein). The IMF orientation may also control the altitude of the magnetic pile-up boundary. Mars Surveyor observations suggest that the altitude of the pile-up boundary rises for eastward IMF orientations but falls for southwest IMF orientations when the subsolar latitude lies in the northern hemisphere (Brain et al. 2005). Without global observations, it is difficult to determine whether this variation results from some as yet unspecific global cause, a local Hall current effect, or local mass loading. The altitude of the magnetic pile-up boundary increases in southern summer, when the stronger crustal magnetic fields in the southern hemisphere approach the subsolar point. Crustal anomalies may also determine magnetosheath densities. Ma et al. (2002) reported simulation results indicating that crustal magnetic features do not cause major distortions of the bow shock, but do have effects on the magnetosheath and the altitude of the ionopause. Simulation results reported by Harnett and Winglee (2003, 2005) predict that mini-magnetospheres extend beyond and replace the magnetic pile up boundary in the presence of crustal anomalies. In the absence of reconnection, IMF draping over strong southern magnetic anomalies at Mars should enhance flank magnetosheath densities by more than a factor of 2 outside dawnward or duskward facing anomalies. In the presence of reconnection, enhancements are far smaller and densities fall within void regions that lie just downstream from the anomalies. According to Vignes et al. (2000) and Bertucci et al. (2005), the bow shock and magnetic pile-up boundaries move inward and outward together, but neither boundary exhibits much response to variations in solar EUV. However, Brain et al. (2005) found evidence for the pile-up boundary moving towards the planet during periods of enhanced solar EUV.

Russell et al. (2006) presented the schematic illustration of the solar wind’s interaction with Venus shown in Fig. 12 and enumerated the plasma structures seen by Pioneer Venus Orbiter (PVO, Colin 1980) at Venus. As in the case of Mars, the interaction results in the formation of a foreshock upstream from the quasi-parallel bow shock, a magnetosheath in which the plasma flow and magnetic field pick up planetary ions, a magnetic barrier with a mixture of solar wind and ionospheric plasmas at the inner edge of the magnetosheath, a (generally) field-free ionosphere, and a mass-loaded magnetotail. Because gradients in the densities of ions with solar wind origins mark each of these boundaries (e.g., Terada et al. 2009), they are imageable in soft X-rays. Table 3 summarizes reported effects of the solar wind pressure/Mach number, IMF direction, and solar EUV on plasma structures at Venus. The ionopause at the inner edge of the magnetic pile up boundary rises from altitudes of 300 to 1000 km as the solar wind dynamic pressure diminishes from 4 to 0.5 nPa (Brace et al. 1982). Distances to the bow shock and the width of the magnetosheath depend primarily on the IMF orientation, the solar wind Mach number, and exospheric neutral densities, rather than solar wind dynamic pressure (Martinecz et al. 2008; Russell et al. 1988; Zhang et al. 2004). Note, however, that Martinecz et al. (2009) found no relation between the location of the terminator bow shock and solar EUV. In the absence of a significant crustal magnetic field, there are no seasonal or diurnal effects.

Fig. 12
figure 12

Schematic illustration of the solar wind interaction with Venus. Solar EUV radiation ionizes the neutral upper atmosphere of Venus. The electron ion thermal pressures suffice to stand off the supersonic solar wind and form a shock. Neutrals formed in the flowing solar wind are carried away by the wind (Russell et al. 2006)

Table 3 Drivers affecting the variability of plasma boundary locations at Venus (Phillips et al. 1985; Knudsen et al. 1987; Russell et al. 1988; Brace et al. 1990)

Slavin et al. (1980) and Tatrallyay et al. (1983) showed that the bow shock at Venus flares more than might be expected based on gasdynamic models, suggesting that mass loading plays an important role. Consistent with this hypothesis, Alexander and Russell (1985) showed that the terminator bow shock moves outward during solar maximum when exospheric neutral densities should be enhanced. Alexander et al. (1986) demonstrated that the bow shock moves away from the planet as the IMF rotates from orientations parallel to the solar wind flow to orientations perpendicular to that flow, i.e., from orientations that do not favor ion pick up to orientations that do. The effect is far greater during solar maximum than solar minimum. As in the case of Mars, the bow shock moves outward for low Mach numbers (Russell et al. 1988). The Venusian bow shock is not circular within the terminator plane, but rather lies further from the planet in the direction perpendicular to the IMF orientation in the plane transverse to the solar wind flow direction, particularly in the direction with the outward pointing electric field where pick up ion effects are expected. Also consistent with the pick up ion effect, this asymmetry in bow shock locations becomes more pronounced during solar maximum. Finally, Zhang et al. (1990) used observed and estimated bow shock locations to infer that the effective radius of the Venusian obstacle to the solar wind lies below the distance to the subsolar ionopause during solar minimum, i.e., that there is a more direct interaction of solar wind plasma with this planet’s ionosphere and exosphere during solar minimum than solar maximum.

The locations of the bow shock at Venus can be used to determine the best value for the polytropic index in the solar wind. Tatrallyay et al. (1984) discussed the strength of the Venusian bow shock by determining magnetic field strength compressions across the bow shock as a function of solar wind Mach number, concluding that polytropic index \(\gamma=1.85\) works best at Venus. The strength of the compressions, and the index \(\gamma\), increase with magnetosonic Mach number and cone angle. The distance to the terminator shock diminishes with magnetosonic Mach number.

Hybrid code simulations predict the principle features of the solar wind’s interaction with unmagnetized planets including the locations of the bow shock and an ion composition boundary between plasma of solar and planetary origin (Bößwetter et al. 2007), but have also made some interesting predictions for the solar wind interactions with Venus and Mars. Whereas simulations predict that the bow shock lies further from the planet in the hemisphere where the convection electric field points inwards towards the planet (e.g., Modolo et al. 2006; Brecht and Ledvina 2006), observations indicate that the bow shock lies further from the planet in the hemisphere containing accelerated pick-up ions, i.e., the hemisphere where the electric field points outward away from the planet (e.g., Dubinin et al. 1996, 1998; Vignes et al. 2002). Shimazu (2001a) reported that the sense of bow shock asymmetries could be reconciled with observations by including the effects of charge exchange in the models. In this case, heavy ions replace the flow of solar wind ions in the magnetosheath. Simulations also predict the occurrence of multiple shock waves (Moore et al. 1991; Shimazu 2001b; Modolo et al. 2006), a feature that would be difficult to identify using single point observations from in situ spacecraft. Because the model results reported by Moore et al. (1991) were relatively insensitive to mass-loading, these authors proposed that solar cycle variations in shock locations result from changes in the dimensions of the magnetic pile up boundary and ionopause rather than changes in the rate of ion pick up. Shimazu (2001b) predicted that the presence of the interplanetary magnetic field constrains the planetary plasma boundary (the ionopause) to an elliptical cross-section. Finally, Martinecz et al. (2009) predicted that pronounced density enhancements extend upstream from the quasi-parallel bow shock at the dawn terminator during intervals of spiral interplanetary magnetic fields. The density jump at the dawn terminator bow shock is much less than that on the dusk side. Brain et al. (2010b) compare the differing predictions of MHD and hybrid models for the solar wind interaction with Mars.

Transient plasma and magnetic field structures are common at both Mars and Venus. Some of these structures might be produced by external solar wind/foreshock drivers, while others result from instabilities at internal plasma boundaries. Within the former category, Collinson et al. (2014) suggested that the significant transient density structures generated by kinetic processes within the Venusian foreshock might drive large amplitude waves on the ionopause and magnetic pile-up boundary of that planet. The same might also be true at Mars, where hot diamagnetic cavities and flow anomalies are also present within the foreshock (Øieroset et al. 2001; Collinson et al. 2015).

However, instabilities may generate other structures. Observations indicate not only the common occurrence of magnetic flux ropes embedded in the ionosphere (Russell and Elphic 1979) but also the frequent occurrence of wavelike structures at the ionopause and clouds of ionospheric plasma above it (Brace et al. 1980, 1982; Acuna et al. 1998). Slowly moving flux ropes are common in the Venusian ionosphere during periods of low solar wind dynamic pressure at solar maximum (Russell et al. 2006), and they are present, albeit much rarer, at Mars (Vignes et al. 2004). According to some estimates these transient events may play a major role in removing planetary plasma (Brace et al. 1982; Russell et al. 1982b; Terada et al. 2002).

The clouds might be caused by a sling shot effect pulling draped magnetic field lines over the planetary poles (Russell et al. 1982b), reconnection and the formation of flux ropes in the magnetosheath (Dreher et al. 1995), or a tearing off of blobs during the final stage of a non-linear Kelvin-Helmholtz instability at the ionopause (Wolff et al. 1980; Gunell et al. 2008). Seeking to confirm the slingshot effect model, Ong et al. (1991) found that clouds are more common on Pioneer Venus Orbiter periapsis passes during which the orientation of the upstream magnetic field changes abruptly, indicating a need for an additional mechanism.

One such mechanism is magnetic reconnection, which should produce magnetic flux ropes in regions of sheared draped magnetosheath magnetic fields outside the ionopause (Dreher et al. 1995). Sheared fields are indeed a natural consequence of rotations in the upstream IMF orientation. However, at Mars there is another way that reconnection can generate flux ropes. Ma et al. (2002) reported simulation results indicating that reconnection of IMF and crustal magnetic field generates mini magnetocylinders of closed magnetic field lines within the Martian magnetosheath, while Harnett (2009) reported simulation results indicating that these cylinders are rapidly dissipating flux ropes with sizes that increase slightly and are on the order of half a planetary radius in the Martian magnetosheath. Brain et al. (2010a) discuss observations indicating that interactions with the solar wind stretch and pinch off loops of crustal magnetic field, resulting in antisunward moving flux ropes filled with ionospheric plasma as shown in Fig. 13. Ropes with greater than 100 nT magnetic field strengths were seen in 1% of Mars Global Surveyor orbits and their estimated diameters were on the order of 2250 km, large compared to the radius of Mars. Consequently, Albee et al. (2001) argued that the ropes might account for up to 5–10% of ion loss at Mars.

Fig. 13
figure 13

Detaching a magnetic flux rope from Mars (Brain et al. 2010a). Panel (a) shows crustal magnetic field lines that are still attached to the planet, but have been stretched tailward by the solar wind. Panel (b) shows a detached loop of crustal magnetic field carrying plasma away from Mars. The dashed line shows the sunward motion of a spacecraft

Reconnection is not the only mechanism for generating flux ropes. Although the subsolar ionopause is generally thought to be stable to both the Kelvin-Helmholtz and flute instabilities (Elphic and Ershkovich 1984), large corrugations are present even here (Russell et al. 1987), leading to a suggestion that they are produced by curvature forces pulling flux tubes draped over the ionopause into the ionosphere. Further from the subsolar ionopause, waves and flux ropes might be produced by the Kelvin-Helmholtz instability (Wolff et al. 1980). Bertucci et al. (2005) inferred the presence of ripples on the Martian magnetic pile up boundary at large solar zenith angles from discrepancies between model normals and those determined from minimum variance analysis. Although sharp ionopause density gradients are the norm at Venus, Duru et al. (2009) noted that they were only observed in 18% of the samples studied at Mars. Noting past work indicating highly fluctuating electron densities in the Martian ionosphere, Duru et al. (2009) attributed the infrequent occurrence of strong ionopause density gradients at Mars to time or spatially dependent phenomena, perhaps the Kelvin-Helmholtz instability at the ionopause. Pope et al. (2009) inferred the presence of giant vortices capable of redistributing and causing the substantial loss of ionospheric plasma at Venus.

Simulations of the Kelvin-Helmholtz instability at the ionopauses of the unmagnetized planets often reach conflicting conclusions. Terada et al. (2002) reported results from a two-dimensional hybrid code simulation which indicate preferential wave growth beginning even at the subsolar point and continuing further antisunward in the hemisphere with an electric field pointing away from the planet, as shown in Fig. 14. By contrast, the waves began to develop only at greater solar zenith angles in the opposite hemisphere. Penz et al. (2004) used an MHD simulation to study the case where flows lie transverse to draped magnetosheath magnetic field lines. Under these conditions, the magnetic field plays no role in stabilizing the instability. The subsolar magnetopause is stable and the non-linear instability develops on the equatorial flanks. They estimated that the atmospheric loss via the Kelvin-Helmholtz instability is comparable to that by other non-thermal loss mechanisms. Amerstorfer et al. (2007) reported results from an MHD simulation with a similar magnetic field configuration, this time in the presence of strong radial gradients in density and velocity at the terminator ionopause. High magnetosonic Mach numbers (increasing compressibility) and greater magnetosheath to ionospheric density ratios diminish the likelihood of the instability. The wavelengths of the fastest growing mode diminish as the density ratios increase. MHD simulation results indicating a sharper ionopause and a greater shear flow led Bößwetter et al. (2007) to conclude that the Kelvin-Helmholtz instability is more likely in the hemisphere in which the convection electric field points towards the planet. Amerstorfer et al. (2010) then considered the evolution of the instability from a linear phase, through a nonlinear phase with regular structures through a turbulent phase with nonlinear structures. They concluded that the instability could account for the atmospheric loss rate estimated from observations. Finally, Möstl et al. (2011) used an MHD simulation to argue that conditions generally do not favor the occurrence of the Kelvin-Helmholtz instability at Venus, with the possible exception of the dayside induced “magnetopause”, or upper boundary of the magnetic barrier, during solar maximum, at dayside locations away from the subsolar point, where magnetosheath flows lie transverse to draped magnetosheath magnetic field lines.

Fig. 14
figure 14

Results from a global multifluid hybrid code simulation for the solar wind interaction with Venus as a function of solar zenith angle (SZA) and altitude (Terada et al. 2002). Panel (a) shows the distribution of magnetic field strengths (colors) while vectors show flow directions for plasma of solar wind origin. Panel (b) shows the densities and flow directions for \(\mbox{O}^{+}\) ions of planetary origin. The ionopause exhibits much greater corrugation in the hemisphere with an upward electric field than the hemisphere with a downward electric field

Global images have the potential to play a decisive role in testing the often conflicting predictions of the various solar wind-planetary interaction mechanisms proposed to occur at Mars and Venus. They can be used to determine occurrence patterns and extent as functions of solar wind and solar cycle conditions, and to quantify the importance of each mechanism to the loss of planetary atmospheres. As an example, take the case of the bow shock asymmetries expected in response to atmospheric loss via ion pick-up. Observations can first be used to test conflicting model predictions indicating that the bow shock lies further from (closer to) the planet in the hemisphere with outward pointing convection electric fields. The degree to which the bow shock is asymmetric can provide information concerning the significance of ion pick-up over time scales ranging from minutes to solar cycles. Similarly, observations of the size, extent, and number of the wavy density structures generated by the Kelvin-Helmholtz instability or with isolated structures associated with the flux ropes generated by magnetic reconnection can be used to determine the importance of these mechanisms as a function of simultaneously measured solar wind conditions.

2.4 The Moon

Despite its tenuous nature, the lunar exosphere remains high on the list of planetary science targets thanks to its complexity and role as an accessible representative of airless bodies in the solar system and the possible presence of water and other potential resources. The role of volatiles in the lunar exosphere is particularly important.

The solar wind and meteoroids deliver protons and other species to the lunar surface at local rates that depend on surface composition, impinging local topography, and the presence of structures such as magnetic anomalies. Solar wind ions weather the surface by creating defects in the lattice that weaken the solid state structure. Because the lunar surface is generally saturated with these volatiles, the implanted species escape the surface and form the volatile lunar exosphere through a variety of processes including sputtering, recoil, and diffusion. These processes deposit H and other volatiles into cold traps and form OH (and possibly water) through chemical alteration of oxygen-bearing minerals. Exospheric volatiles are reclaimed by the solar wind as picked-up photoions and charge-exchange products. Global imaging of the total lunar exosphere including all species at regional scales as functions of solar zenith angle and the plasma and space environment will lead to a unified understanding of the plasma, exospheric, and geologic Moon.

The Lunar Atmosphere Dust and Environment Explorer (LADEE, Elphic et al. 2014) Neutral Mass Spectrometer (NMS, Mahaffy et al. 2014) confirmed the presence of water in the equatorial lunar exosphere for brief periods early in the instrument turn on/warm up period. These detections of non-polar exospheric water occurred preferentially near the radiant of episodic meteor stream encounters. Water densities of \(2\mbox{--}3\times10^{8}~\mbox{m}^{-3}\) during these meteor shower events (Benna et al. 2015a) correlate nicely with LADEE Lunar Dust Experiment (LDEX, Horányi et al. 2014) dust stream occurrences.

Although LADEE established that lunar volatile gases like water can be released by the impact of solar system objects like meteoroids in the equatorial region, volatiles can also be released from the interior of the Moon, through moonquakes (Cook and Stern 2014). Additionally, they can be synthesized in the upper layer of the lunar regolith by the solar wind. Once released, they are transported across the lunar surface until they either escape to space or become trapped in cold permanently shadowed regions (PSRs) that have maintained temperatures below 100 K for billions of years. Of particular interest is water trapped in these PSRs.

In fact, there has been observational verification of an active water and hydroxyl environment (i.e., water cycle) at the Moon including Lunar Crater Observation and Sensing Satellite (LCROSS, Schultz et al. 2010) confirmation of water existing within the lunar polar cold traps (Colaprete et al. 2010; Schultz et al. 2010). Other evidence includes data from a set of IR sensors showing an OH veneer that extends all the way down to the lunar equator, and which may even possess a present-day, dynamic diurnal component (Pieters et al. 2009; Clarke et al. 2009; Sunshine et al. 2009). The distribution of water in the lunar polar regions is heterogeneous on all observed scales (Mitrofanov et al. 2010).

However prior to deposition into cold traps, the volatiles must be transported some distance across the lunar surface. Volatile mobility depends on many parameters including species, surface composition, and temperature. For example, the argon density distribution results from a surface interaction, an excess of adsorption over desorption on the nightside as the lunar surface cools, so its density peaks at the terminator where the surface heats up (Hodges 1977). Helium, on the other hand, is not adsorbed onto the surface so it spends more time on the cold nightside than on the warmer dayside because the lateral extent of its trajectories is proportional to temperature (Hodges 1973, 1975). Consequently, He density peaks on the nightside. Of course, the scale height and its dependence on temperature also play a role.

In general, the cold nightside lunar atmosphere is dominated by non-condensible species, including He, detected by Apollo-era instrumentation, and Ne and H2, as observed by LADEE and the Lunar Reconnaissance Orbiter (LRO, Tooley et al. 2010). LADEE also confirmed the presence of argon at the equator and the Lyman-Alpha Mapping Project (LAMP, Gladstone et al. 2010) placed limits on Ar at the poles (Hodges and Mahaffy 2016; Grava et al. 2015; Benna et al. 2015b). These in situ observations, when coupled with global data on the structure of the lunar exosphere including local time dependence and vertical scale heights, are essential for determining production rates and polar cold trapping efficiencies. With guidance from modeling efforts, global images of total exospheric content could determine the constituents of the lunar atmosphere over the poles.

Global imaging will also reveal the relationship between the time-variable solar wind flux and composition and the lunar exosphere. This would be accomplished in a manner similar to what LRO did in situ (Feldman et al. 2012) by showing that the surface He density exhibits variations responding to changes in the solar wind alpha flux (see also Benna et al. 2015b). Of course, global imaging would provide an overall perspective on this process not possible with local in situ measurements. Global images of the lunar atmosphere can also be used to study the behavior of the lunar exosphere as the Moon moves in and out of the terrestrial magnetotail, modulating solar wind sputtering and enabling identification of the distant terrestrial magnetopause and possibly bow shock. Imaging can also reveal the global effects of meteoroid bombardment (e.g., Colaprete et al. 2016).

Soft X-ray observations of the lunar exosphere complement and validate model predictions for the dominant contributors to the exospheric column density. Furthermore, because soft X-ray imaging relies on the presence of the solar wind, global imaging will reveal the shape and extent of structures effected by the solar wind-lunar interaction. Plasma structures in the vicinity of the Moon include a low density wake (Lyon et al. 1967; Zhang et al. 2014) and mini-magnetospheres above magnetic anomalies (Wieser et al. 2010) that could be imaged globally using soft X-ray emission. In addition to morphology, global imaging will reveal aspects of the interaction that can be quantitatively compared to model predictions, for example the extent over which solar wind ions impact the lunar surface beyond the terminator (Collier et al. 2014).

Global images will supply the key to a unified understanding of the plasma, exospheric, and geologic Moon. They will provide information on the exospheric content as a function of altitude and location above the lunar surface that can be correlated to geologic regions. Global imaging will also reveal properties of the solar wind plasma-lunar interaction, such as wake morphology and how magnetic anomalies affect solar wind implantation.

3 Soft X-Ray Intensities from Solar Wind Charge Exchange

Solar wind charge exchange is responsible for EUV and soft X-ray emission not only in regions of the solar system where the solar wind interacts with neutral gases from objects such as comets, and planetary exospheres such as Earth’s geocorona, but also with the interstellar gas flowing through the heliosphere (see Sect. 4.2). Charge exchange leaves the product ion in an excited state which then returns to the ground state through the emission of one or more photons. Since the bulk of the ions in the solar wind are highly ionized, most of these photons are in the soft X-ray and extreme ultraviolet. As shown in Fig. 15, the resulting spectrum is extremely rich. However, this spectrum is generally observed at relatively low spectral resolution so that the bulk of the lines are severely blended. This section reviews the production of EUV and soft X-rays by solar wind charge exchange. In particular it describes the many factors required to determine the spectrum seen by an observer looking along a single LOS. Calculating the spectrum involves many quantities that are poorly known. However, we will show that working at lower resolution makes the problem in some ways more tractable.

Fig. 15
figure 15

Model solar wind charge exchange spectrum (similar to Fig. 1 from Koutroumpa et al. 2009b)

The intensity, \(I\) (in photon \(\mbox{cm}^{-2}\,\mbox{s}^{-1}\)), of the emissions from transition \(j\), for species \(s\), in charge state \(q\), seen by an observer looking along a LOS is given by the integral through the emitting region(s) along that line of sight:

$$ I_{j} = \int P_{sqj} dl = \sum_{n} \int n_{n} n_{q} v_{rel} \sigma_{sqn} b_{sqj} d\varOmega dl/4\pi, $$

where \(P_{sqj}\) is the volume emission rate (photons \(\mbox{cm}^{-3}\,\mbox{s}^{-1}\)) for a specific transition \(j\) (with photon energy \(E_{j}\)) from a specific charge exchange collision of the solar wind ion species (denoted \(s\)) in charge state \(q\) with a neutral target \(n\). The summation over \(n\) reflects the reality that there may be multiple neutral species, though in many cases we need only consider H and He (for the diffuse heliospheric emission) or only H (for the Earth’s magnetospheric emission). The relevant charge exchange cross section at the appropriate collision energy is \(\sigma_{sqn}\), and \(b_{sqj}\) is a branching ratio in the product ion species for the transition of interest. The branching ratio is the fraction of ions undergoing charge exchange between \(n\) and \(sq\) that relax through transition \(j\). The densities \(n_{n}\) and \(n_{q}\) are those of the neutrals and ions respectively, while \(v_{rel}\) is the relative velocity of the neutrals and ions:

$$ v_{rel}\sim\bigl(v_{r}^{2}+v_{therm}^{2} \bigr)^{\frac{1}{2}} $$

where \(v_{r}\) is the bulk velocity of the ions, it being supposed that the thermal velocity of the target neutrals is small. The quantity \(v_{therm}\) is \(3kT/m_{p}\) for the solar wind ions. The photon flux within some field of view is the integral of intensity over solid angle increment \(d\varOmega\). This equation is generally rewritten as:

$$ I_{j} = \int P_{sqj} dl = \int n_{n} n_{p} v_{rel} \frac{n_{q} n_{s}}{n_{s} n_{p}} \sigma_{sqn} b_{sqj} d\varOmega dl/4\pi, $$

where \(n_{p}\) is the solar wind proton density. To calculate the integral intensities along specific lines of sight, each of these factors must be considered in detail.

It should be noted that in many cases, instrumental resolution is insufficient to isolate individual lines. For a bandpass containing emission from many different transitions (\(j\)) from many ion species (\(sq\)) charge exchanging with different neutral targets (\(n\))

$$ I= \sum_{j} \sum_{sq} \sum_{n} \int n_{n} n_{p} v_{rel} \frac{n_{q} n_{s}}{n_{s} n_{p}} \sigma_{sqn} b_{sqj} d\varOmega dl/4\pi $$

which can be rewritten as

$$ I = d\varOmega/4\pi\sum_{n} \int n_{n} n_{p} v_{rel} dl \biggl[\sum _{j} \sum_{sq} \frac{n_{q} n_{s}}{n_{s} n_{p}} \sigma_{sqn} b_{sqj} \biggr] \equiv\sum _{n} \frac{d\varOmega}{4\pi} Q_{n} \varsigma_{n} $$


$$ Q_{n} \equiv \int n_{n} n_{p} v_{rel} dl\quad\mbox{and}\quad \varsigma_{n} \equiv\sum_{j} \sum _{sq} \frac{n_{q} n_{s}}{n_{s} n_{p}} \sigma_{sqn} b_{sqj}. $$

Here we have assumed that \(n_{q}/n_{s}\), \(n_{s}/n_{p}\), and \(\sigma_{sqn}\) (which must be a function of \(v_{rel}\)) are at least relatively constant along the line of sight. This formulation segregates the bulk properties of the solar wind and its neutral targets from the atomic data.

An alternate formulation seen in the literature for both individual lines and band passes is more convenient for calculating energy fluxes (in \(\mbox{eV}\,\mbox{cm}^{-2}\,\mbox{s}^{-1}\)):

$$ F_{j} = \int E_{j}P_{sqj}dl = Q \alpha_{j} $$


$$ \alpha_{j} \equiv E_{j} \frac{n_{q} n_{s}}{n_{s} n_{p}} \sigma_{sqn} b_{sqj} $$

that is, the energy-weighted cross section, which is also called an emission scale factor.

As we will see below, \(Q\) can be derived from models with some reasonable degree of confidence, while \(\varsigma\), a production factor, requires atomic data that is, in many cases, unknown. In Sect. 3.1.3 we will demonstrate that \(\varsigma\) has been derived from observations for the broad Röntgensatellit (ROSAT, Trümper 1992) \(\frac{1}{4}~\mbox{keV}\) bandpass, which allows simulations of the entire magnetosheath in that and similar bands.

3.1 Theoretical and Observation-Inferred Charge-Exchange Cross Sections

3.1.1 Predicted Charge-Exchange Cross-Sections

The collision of an ion with a neutral target can result in the transfer of an electron from the neutral atom or molecule to the ion, i.e., charge exchange. The incident ion with charge, \(q\), is represented as \({\mathrm{M}}^{q+}\), where M is a minor species such as O, N, C, and Fe. Of greatest interest are high charge state ion species such as \(\mbox{O}^{7+}\) or \(\mbox{C}^{6+}\) that are abundant in the solar wind and produce X-rays upon recombination. The target neutral species is designated \(\mbox{B}\), where \(\mbox{B}\) can be H2O for comets, H for the Earth’s exosphere, H and He for interplanetary space, or other neutral species as required. The relevant charge exchange reaction can be written:

$$ {\mathrm{M}}^{q+} + \mbox{B} \rightarrow{\mathrm{M}}^{(q-1)+} + \mbox{B}^{+} $$

A key example of this reaction in the terrestrial exosphere or in the heliosphere is:

$$ \mbox{O}^{7+} + \mbox{H} \rightarrow \mbox{O}^{6+} + \mbox{H}^{+} $$

For highly ionized recipients (i.e., large values of \(q\)), the product ion species (i.e., \({\mathrm{M}}^{(q-1)+}\)) is invariably left in a highly excited state such that a radiative cascade follows the collision as the excited ion de-excites to the ground-state. For large values of \(q\) at least one of the photons in this cascade is an EUV or a soft X-ray photon.

At solar wind energies (i.e., \(\sim1~\mbox{keV/amu}\)), the cross section \(\sigma_{sqn}\) for this type of charge exchange collision is very large, much greater than geometrical based on the dimensions of the interacting particles. Figure 16 shows interaction potential energy curves versus inter-nuclear separation for two atomic species undergoing a charge exchange collision (Isler 1994). The dashed line indicates results for incident reactants \(\mbox{O}^{8+}\) and H, which experience a weak point charge-induced dipole interaction whose energies remain almost constant at larger distances. By contrast, the outgoing reaction products \(\mbox{O}^{7+}\) and \(\mbox{H}^{+}\) experience a strong Coulomb interaction. The curves include the energies of the hydrogen-like \(\mbox{O}^{7+}\) excited to states with principle quantum number \(n\). If there is a suitable resonance for the electron in the system, then the charge exchange reaction can be said to take place with some probability at a curve crossing. The curve crossings for high values of \(n\) take place at large radii (e.g., \(r\approx9 a_{0}\) for \(n\approx5\), where \(a_{0}=0.54\times10^{-8}~\mbox{cm}\) is the Bohr radius). Hence, for a reaction probability of about 0.5 the cross section would be \(\sigma\approx\pi\times(9 a_{0})^{2} \approx5 \times10^{-15}~\mbox{cm}^{2}\).

Fig. 16
figure 16

Interaction of potential energy curves versus inter-nuclear separation for two atomic species undergoing a charge exchange collision. (Adapted from Isler 1994, see text for additional information)

The probability of charge-exchange depends not only on the principal quantum number, but on all of the values that describe the state into which the electron is initially inserted. The cross-sections, \(\sigma _{sqn}\), referred to above are the total cross-section over all possible initial states, while the details depending on the initial states are hidden in the \(b_{sqj}\). Calculated cross-sections must be constructed from those of each initial state. Total cross-sections can measured, but \(n,l,m\) resolved cross-sections produce far more insight into the physical process.

Figure 17 demonstrates the resonance process that necessitates the energy of the final state of the recipient ion product in the charge exchange reaction. It shows the electron potential energy as function of distance along the inter-nuclear axis for (in this case) the \(\mbox{B}^{5+} + \mbox{H} \rightarrow\mbox{B}^{4+}(n=3) + \mbox{H}^{+}\) system at a time during the reaction when the nuclei are 9 atomic units (a.u.) apart (Cravens 2002). Here 1 a.u. of distance is 1 a0 and 1 a.u. of energy is 1 Rydberg or 27.2 eV. This internuclear distance is a favorable one because the electron energy can remain about the same (the resonance) and “move over” from the region near the \(\mbox{H}^{+}\) nucleus over a low energy barrier to the region near the \(\mbox{B}^{x+}\) nucleus.

Fig. 17
figure 17

Electron potential energy as a function of distance along the inter-nuclear axis for the \(\mbox{B}^{5+} + H \rightarrow \mbox{B}^{4+}(n=3) + \mbox{H}^{+}\) system at a time during the reaction when the nuclei are 9 atomic units (a.u.) apart (adapted from Cravens 2002)

The simple classical over-barrier (COB) collision model provides approximate cross sections and excitation levels (Mann et al. 1981). The cross sections are energy-independent and apply only to relatively low collision energies but are appropriate for solar wind ions. For a fully stripped recipient species \(\mathrm{M}^{q+}\) colliding with a neutral target species B that has an ionization potential of \(I_{p}\), the energy defect is \(\Delta{E} = q^{2}/2n^{2} - I_{p}\) and the curve crossing distance is \(R_{x} \approx(q-1)/\Delta{E}\), where \(I_{p} = 0.5~\mbox{a.u.}\) for atomic H. The COB cross section is then \(\sigma\approx\pi{R}^{2}_{x}\) and the most likely excitation level (i.e., the principle quantum number) is given by \(n \le q\{2 I_{p} [1 + (q-1)/(2q^{1\over 2} + 1)]\}^{-{1\over 2}}\). For the \(\mbox{O}^{8+} + \mbox{H}\) collision, this gives \(n \approx5\) for the product \(\mbox{O}^{7+}\) and \(\sigma\approx250\ a^{2}_{0} \approx7\times10^{-15}~\mbox{cm}^{2}\).

3.1.2 Experimental Measurements of Charge-Exchange Cross Sections

Numerous laboratory measurements of high charge state ion collisions with neutrals have been made over the years. Gilbody (1986) reviewed some of the earlier experimental work, reporting for example on laboratory measurements and theoretical calculations for \(\mbox{C}^{6+} + \mbox{H}\) charge exchange cross sections as a function of energy. In this case the measured charge exchange cross section is \(3\times10^{-15}~\mbox{cm}^{2}\) for a collision energy of 1 keV/amu, which greatly exceeds the geometrical cross section. Janev and Winter (1985), and Janev et al. (1983, 1988) reported measurements of state-selective cross sections that indicated the product ion is left excited with a high principle quantum number (e.g., \(n\approx4\)).

The cross sections \(\sigma_{sqn}\) are velocity dependent, but approximately constant for most solar wind species as a function of the relative velocity between the interacting particles over reasonable velocity ranges. While some important emission lines like \(\mbox{O}^{6+}\) and \(\mbox{O}^{7+}\) have been reasonably well characterized, many others that contribute to the emissions with energies \(<500~\mbox{eV}\) that are relevant to imaging solar wind interactions with the planetary objects are not (see Smith et al. 2014, for a discussion of alternative methods).

Beiersdorfer et al. (1999, 2000, 2001, 2003), Wargelin et al. (2008), Greenwood et al. (2001), Mawhorter et al. (2007), and Betancourt-Martinez et al. (2014) present some relatively recent experimental measurements of X-ray emissions generated by charge exchange. (See also the review within Krasnopolsky et al. 2004.) The more recent experiments include a wide variety of target (i.e., H2O and CO2) and incident ion species and charge states. A variety of experimental methods and incident energies were employed. For example, Greenwood et al. (2001) measured the X-ray spectrum emitted during the charge exchange process using a crossed-beam experiment in addition to determining the initial and final charge states of the recipient ions.

Beiersdorfer et al. (2003) used a microcalorimeter to measure the X-ray spectrum generated when trapped ions interact with neutral targets in the Lawrence Livermore electron beam ion trap (EBIT-I). For \(\mbox{O}^{7+}\) ions interacting with neutrals like CO2, the resulting helium-like emission spectrum from \(\mbox{O}^{6+}\) produces X-ray transitions such as those observed at comets with especially strong lines near 570 eV. In particular, the microcalorimeter detected strong emission from the forbidden transition 1s2s \({}^{3}\mbox{S}_{1}\)\(\mbox{1s}^{2}\) \({}^{1}\mbox{S}_{0}\) at 564 eV. Figure 18 shows an energy level diagram for \(\mbox{O}^{6+}\) illustrating the forbidden, resonant, and inter-combination lines for the \(n=2\) to \(n=1\) transition.

Fig. 18
figure 18

Energy level diagram for \(\mbox{O}^{6+}\) (\(\mbox{O}^{6+}\) X-ray emission lines)

Employing ACE measurements of the solar wind composition, Whittaker et al. (2016) calculated the charge exchange emission scale factors for \(\mbox{O}^{7+}\) and \(\mbox{O}^{8+}\). They showed that the scale factors peak sharply near \(8.2\times10^{-17}~\mbox{eV}\,\mbox{cm}^{2}\) at solar wind velocities of \(\sim400~\mbox{km}\,\mbox{s}^{-1}\), diminishing rapidly for lower velocities and more gradually for higher velocities.

Values used for the emission scale factor for the X-ray band with \(E>50~\mbox{eV}\) range from \(6\times10^{-16}~\mbox{eV}\,\mbox{cm}^{2}\) (Cravens et al. 2001; Robertson and Cravens 2003b) to \(9.4\times 10^{-16}~\mbox{eV}\,\mbox{cm}^{2}\) (Pepino et al. 2004) to \(1.5\times 10^{-15}~\mbox{eV}\,\mbox{cm}^{2}\) (Robertson and Cravens 2003b; Cravens et al. 2001) for the slow solar wind (all of whom cite Schwadron and Cravens 2000 in one way or another for their values) and \(3.3\times10^{-16}~\mbox{eV}\,\mbox{cm}^{2}\) for the fast solar wind (Pepino et al. 2004).

Summarizing, the charge exchange cross sections for many interactions that generate emissions in the \(0.1\mbox{--}0.284~\mbox{keV}\) band remain both poorly understood and poorly determined. Theory and observations indicate a wide range of values for charge exchange cross-sections and their corresponding emission scale factors for H interacting with \(\mbox{O}^{7+}\) and \(\mbox{O}^{8+}\). As we shall see (Sects. 5 and 6), the band-integratedth cross-section for solar wind charge exchange is best estimated from well calibrated observations in space.

3.1.3 A Band-Averaged Production Factor for the 1/4 keV Band

As will be described more fully in Sect. 6.1, \(\frac{1}{4}~\mbox{keV}\) X-ray emission from the magnetosphere was observed by ROSAT as a temporally variable background component which was measured and removed from the ROSAT All-Sky Survey. The temporally variable emission was later shown to be well correlated with the solar wind flux. The relation between observed X-ray emission and the solar wind flux is

$$ \mbox{ROSAT counts s}^{-1}\,\mbox{degree}^{-2} =(0.083 \pm2.26)\times10^{-2}+(0.186\pm0.009) \biggl[\frac {n_{sw}v_{sw}}{10^{8}}\biggr] $$

where \(n_{sw}v_{sw}\) is in \(\mbox{cm}^{-2}\,\mbox{s}^{-1}\) (Kuntz et al. 2015). Given the \(Q \equiv\int n_{n} n_{p} v_{rel} dl\) through which ROSAT observed, one can calculate the \(\frac{1}{4}~\mbox{keV}\) band-averaged production factor. Modeling the magnetosheath for each observation of the All-Sky Survey is computationally prohibitive. Kuntz et al. (2015) used a suite of extant MHD runs to do the equivalent; given the ROSAT observing geometry, they determined the typical \(Q\) through which ROSAT would have observed as a function of the solar wind flux.

$$ \biggl[\frac{n_{sw}v_{sw}}{10^{8}} \biggr] = (0.037\pm 0.189)+(20.68+0.41) \biggl[ \frac{Q}{10^{20}} \biggr] $$

where \(Q\) is in \(\mbox{cm}^{-4}\,\mbox{s}^{-1}\). The combination of these two relations yields

$$ \mbox{ROSAT counts s}^{-1}\,\mbox{degree}^{-2} = (3.86 \pm0.20)\times10^{-20}Q. $$

We can then use this production factor to determine the ROSAT \(\frac{1}{4}~\mbox{keV}\) count rate for a given \(Q\) and, given a model for the shape of the spectrum of the emission, can convert this to the count rate for any instrument with a similar band-pass.

Using the available atomic data rather than ROSAT observations, Robertson et al. (2009b) calculated an equivalent production factor for the slow solar wind interacting with neutral H: \(8.51\times10^{-21}\) count \(\mbox{cm}^{4}\,\mbox{deg}^{-2}\) which is a factor of 4.53 smaller than the value derived by Kuntz et al. (2015) above. Koutroumpa et al. (2009a), also using a collection of atomic data, created a solar wind charge exchange spectrum for a slow solar wind interacting with neutral H. The Koutroumpa et al. (2009a) spectrum contains 2.5 more flux than the Robertson et al. (2009b) spectrum, and thus would have a production factor of \(2.13\times10^{-20}\) count \(\mbox{cm}^{4}\,\mbox{deg}^{-2}\), only a factor of 1.79 lower than the Kuntz et al. (2015) measurement. Since the atomic data used is uncertain and likely to be missing many of the fainter transitions, this agreement between model and measurement is surprisingly good.

3.2 The Branching Ratio and Spectra

Once the charge transfer collision has occurred and the product ion is produced with a high principle quantum number, \(n\), and angular momentum quantum number, \(l\), a radiative cascade ensues, subject to the relevant selection rules (e.g., \(\Delta{l}=\pm1\) for dipole-allowed transitions), so that the ion eventually ends up in the ground-state.

The details of this cascade depend on the set of radiative transition probabilities (or Einstein \(A\) coefficients), and are encompassed in the branching ratio coefficient, \(b_{sqj}\), which appears in the intensity expression given above (see Eqs. (1) and (3)). The coefficient in this equation is an average that must also include information on the initial quantum numbers of the product ion. For example, if an \(\mbox{O}^{5+}\) is created in the \(n=4\) and \(l=1\) state (2s4p) by charge exchange of \(\mbox{O}^{6+}\) with a neutral, then it can radiate to the following states: \(\mbox{2s}^{2}\) with fraction 0.77, 2s3s with fraction 0.11, and 2s3d with fraction 0.04. The 2s2p state will then radiate 100% to the ground state \(\mbox{2s}^{2}\). The difficult task of finding \(b_{sqj}\) for all species and charge states relevant to solar wind charge exchange remains only partially and approximately complete.

For the solar wind ions of interest here, some of the resulting transitions are in the EUV and soft X-ray parts of the spectrum. For example, \(\mbox{O}^{7+}\) ions (produced by charge exchange from solar wind \(\mbox{O}^{8+}\) ions) generate a hydrogen-like spectrum (\(\mbox{O}^{7+}\) emission lines in X-ray astronomy notation) while \(\mbox{O}^{6+}\) ions (from \(\mbox{O}^{7+}\)) generate a helium-like spectrum (\(\mbox{O}^{6+}\) emission lines), as discussed earlier. These lines are in the soft X-ray part of the spectrum. Other recipient species with different charge states produce different spectra.

The detailed X-ray spectrum resulting from solar wind charge exchange (SWCX) depends on the abundances (or fluxes) of the highly-charged heavy ions in the solar wind (e.g., \(\mbox{C}^{6+}\), \(\mbox{O}^{7+}\), \(\mbox{O}^{8+}\), \(\mbox{Mg}^{12+}\), \(\mbox{Fe}^{13+}\), etc.). These abundances depend on where on the Sun the solar wind originated, as discussed by von Steiger et al. (2000), and Schwadron and Cravens (2000). For example, the slow (\(300~\mbox{km}\,\mbox{s}^{-1}\)) solar wind originates from a hotter solar corona and has a relatively higher \(\mbox{O}^{7+}/\mbox{O}^{6+}\) ratio than the fast (\(700~\mbox{km}\,\mbox{s}^{-1}\)) solar wind that originates from cooler parts of the solar corona. Solar wind composition is discussed in more detail in the next section.

Given the solar wind heavy ion abundances and relevant charge exchange cross sections and radiative cascade probabilities, both the EUV and soft X-ray spectra and the efficiencies of the SWCX mechanism can be determined. Several authors have undertaken this exercise both for detailed spectra and for broad-band X-ray emission bands with and without instrumental response functions included for specific observations such as those made by ROSAT (cf., Kharchenko and Dalgarno 2001; Pepino et al. 2004; Krasnopolsky et al. 2004; Robertson et al. 2009b). Figure 19 shows a cometary SWCX spectrum calculated by Kharchenko and Dalgarno (2000). The strong \(\mbox{O}^{6+}\) lines near 570 eV are particularly obvious but a large number of other lines are present at lower energies.

Fig. 19
figure 19

The total emission spectra from the single (dashed lines) and sequential (solid lines) collision regions of a cometary atmosphere normalized to unit flux of solar wind (SW) ions. Collisions of \(\mbox{C}^{4+}\), \(\mbox{C}^{5+}\), \(\mbox{C}^{6+}\), \(\mbox{N}^{7+}\), \(\mbox{O}^{6+}\), \(\mbox{O}^{7+}\), \(\mbox{O}^{8+}\), and \(\mbox{Ne}^{8+}\) with cometary neutral atoms and molecular constituents are considered. The energy resolution of the discrete spectral lines \(\varGamma\) is taken arbitrarily as 1 eV (adapted from Kharchenko and Dalgarno 2000)

3.3 The Flux of High Charge State Solar Wind Ions

We take the flux (\(F_{sq}\)) of ions of species \(s\) with charge state \(q\) to be proportional to that of protons, i.e., \(F_{sq} \approx f_{sq} n_{sw} u_{sw}\), where \(f_{sq}\) is the relative abundance of individual heavy ion species \(s\) in charge state \(q\), \(n_{sw}\) is the proton density and \(u_{sw}\) is solar wind proton velocity. Consistent with this assumption, we note that Neugebauer et al. (2000) reported a close correlation between OMNIWeb proton and SOHO O fluxes over a period of 8 days during 1996. Whittaker et al. (2016) found that simulated and observed integrated line-of-sight charge exchange soft X-ray intensities agree better when the ratio of oxygen to proton flux in the model is held constant than when it is allowed to track that observed simultaneously in the solar wind. Figure 20 shows ACE SWICS (Gloeckler et al. 1998) observations of the flux of \(\mbox{O}^{6+}\) ions versus ACE SWEPAM (McComas et al. 1998) observations of the flux of protons from 1998 to 2011. The plot exhibits considerable scatter, but the two fluxes are well correlated, confirming the assumption of a relatively constant ratio.

Fig. 20
figure 20

The SWICS \(\mbox{O}^{+6}\) flux versus the SWEPAM proton flux from the ACE data binned to 1 day intervals covering 1998 to 2011. The blue line is the canonical O/H value of 0.00045. The red boxes are the means for \(\delta\log{(n_{p}v_{p})}=0.1\) bins, while the bars show the dispersion \([\sum(x^{2})/n-(\sum(x)/n)^{2}]^{0.5}\) in each bin

Nevertheless, the relative heavy ion abundances differ for the slow and fast solar wind flows, with the slow components originating at equatorial latitudes and the fast components originating at coronal holes. Table 4 summarizes average slow solar wind proton and He densities, bulk speeds, and thermal speeds at Earth. Near the ecliptic plane, 90% of slow-stream density measurements lie between 1 and \(20~\mbox{cm}^{-3}\). von Steiger et al. (2013) report H/O ratios of 1500, which gives total O densities of \(2.0\times10^{-3}~\mbox{cm}^{-3}\) for proton densities of \(3~\mbox{cm}^{-3}\). The slow solar wind composition is very biased towards elements with low first ionization potentials and matches that of the solar corona (Schwadron et al. 1999; von Steiger et al. 2000). Predominant species are: \(\mbox{C}^{5+}\), \(\mbox{C}^{6+}\), \(\mbox{O}^{6+}\), and \(\mbox{O}^{7+}\) (Schwadron and Cravens 2000; von Steiger et al. 2000).

Table 4 Solar wind properties (Ebert et al. 2009)

Table 4 also summarizes typical fast solar wind proton and He densities, bulk speeds, and thermal speeds. Near the ecliptic plane, 90% of fast-stream H density measurements lie between 1 and \(10~\mbox{cm}^{-3}\). Corresponding average high charge state O densities are slightly greater than those in the slow solar wind, \(\sim2.2\times10^{-3}~\mbox{cm}^{-3}\). The fast solar wind composition exhibits less of a bias towards elements with low first ionization potentials and matches that of the photosphere. Predominant species are: \(\mbox{C}^{5+}\), \(\mbox{N}^{5+}\), \(\mbox{O}^{6+}\), and \(\mbox{Ne}^{8+}\).

As illustrated in Fig. 21, fast solar wind streams overtake slow solar wind streams, creating compressional corotating interplanetary regions (CIRs, Pizzo 1978) with enhanced densities and magnetic field strengths that spiral outward from the Sun. Fast-moving, outward-propagating, interplanetary coronal mass ejections (ICMEs) occasionally disrupt this two-stream pattern, particularly at the peak and during the declining phase of the solar cycle. The azimuthal extents of ICMEs at Earth are about four times greater than their \(\sim0.1~\mbox{AU}\) half-widths (Russell and Mulligan 2002). Typical ICME proton and He densities, bulk speeds, and thermal speeds are again listed in Table 4 (Ebert et al. 2009). Near the ecliptic plane, 90% of ICME density measurements lie between 1 and \(10~\mbox{cm}^{-3}\). The overall chemical composition of ICMEs resembles that of the slow solar wind, but the He density is frequently enhanced, with the ratio of He to proton densities often exceeding 8% and sometimes exceeding 10%. Minor ions with \(Z>2\) in CMEs have enhanced densities with respect to protons. This has been particularly well-documented in the Fe/H ratios (e.g., Ipavich et al. 1986; Mitchell et al. 1983; Bame et al. 1979). Furthermore, the ionization temperatures inferred from the charge state distributions are also frequently, but not always, elevated in comparison to coronal hole flow, which is also high-speed, and interstream flow. For example, coronal hole oxygen charge states indicate a 1.3 MK ionization temperature whereas CME-related charge states indicate an ionization temperature above 2 MK (Galvin 1997). Shocks propagating ahead of the ICMEs compress the local plasma, enhancing densities and magnetic field strengths.

Fig. 21
figure 21

Fast and slow solar wind stream structure (Pizzo 1978)

The limited observations available indicate that the composition of high charge state heavy ion populations in the magnetosheath faithfully reflect those of their parent populations in the solar wind at much higher densities (Gloeckler et al. 1986). In the absence of any proposed process that would preferentially remove or accelerate high charge state heavy ion populations crossing the magnetopause, the same must be true for the ion populations in the LLBL, cusps, and mantle. Whether or not their differing Larmor radii enable ions with different mass to charge state ratios to reach different locations remains unknown.

Finally, the flux of ions \(F_{sq}\) can be written as \(n_{sq}V_{rel}\), where \(n_{sq}\) is the density of species \(s\) with charge state \(q\) and \(V_{rel}\) is the effective velocity. Since charge exchange continues in hot plasmas like the magnetosheath and cusps with significant thermal velocities \(V_{th}\) even when the bulk velocity \(V_{bulk}\) vanishes, \(V_{rel}\) can be defined as \((V^{2}_{th} + V^{2}_{bulk})^{1\over 2}\).

3.4 Neutral Densities in the Outer Exosphere

Earth’s exosphere is primarily H at all radial distances of interest here. It pervades all of the regions mentioned above, including the magnetosphere, cusps, boundary layers, the magnetosheath, and the near-Earth solar wind. Chamberlain (1963) presented a model for the exosphere in which there is a transition with radial distance from atmospheric densities dominated by orbiting particles to densities dominated by ballistic and escaping particles, to one dominated solely by escaping particles. In the latter region, temperatures remain constant with increasing radial distance at values \(\sim12\%\) of those at the critical level deeper within the atmosphere at the altitude where collisions first become negligible. Hodges (1994) presented the results of a Monte Carlo simulation for the seasonal and solar cycle variation of exospheric densities at distances from Earth ranging from the exobase to \(10~\mbox{R}_{E}\) that included atoms on ballistic trajectories and collisions. The model does not include the enhanced hot exospheric neutral densities in the vicinity of the cusps and magnetopause that result from the charge exchange of energetic protons and exospheric hydrogen. Because exospheric temperatures are determined by the relatively cold values at the critical level, the effective temperatures governing the interaction of exospheric neutrals and geospace plasmas are determined almost exclusively by the thermal and bulk velocities of the geospace plasmas.

By contrast to the sharp plasma density gradients predicted by models for the solar wind-magnetosphere interaction, exospheric models predict relatively smooth and gradual transitions in neutral density. Hodges’ exospheric model predicts that neutral densities at \(10~\mbox{R}_{E}\) peak near the equator during the equinoxes, but at off-equatorial latitudes during the solstices. They fall off approximately as \(R^{-3}\), where \(R\) is the radial distance from Earth. Dayside values at \(10~\mbox{R}_{E}\) near local noon increase from \(>22.5~\mbox{cm}^{-3}\) at solar wind radio fluxes at 2800 MHz, \(\mbox{F10.7} = 230\) to \(>37.5~\mbox{cm}^{-3}\) for \(\mbox{F10.7} = 80\) during the equinoxes. Dayside values at \(10~\mbox{R}_{E}\) near local noon increase from \(>27.5\) for \(\mbox{F10.7} = 230\) to \(>37.5~\mbox{cm}^{-3}\) for \(\mbox{F10.7} = 80\) during the solstices. Consistent with the predictions of this theoretical model, Zoennchen et al. (2011) inferred exospheric neutral densities from TWINS observations of scattered solar \(\mbox{Ly}\alpha\) finding a gradual transition from greater nightside (\(\sim45~\mbox{cm}^{-3}\)) than dayside (\(\sim22.5~\mbox{cm}^{-3}\)) neutral densities at \(10~\mbox{R}_{E}\) from Earth during solar minimum. Nevertheless, they noted the possibility of 100% errors in the model at these distances from Earth. Figure 22 compares their results with those from previous empirical studies. Finally, Fuselier et al. (2010) presented a case study of energetic neutral atom observations from which they inferred a neutral density of only \(8~\mbox{cm}^{-3}\) at the subsolar point. Given the uncertainties inherent in both the models and observations, neutral densities at \(10~\mbox{R}_{E}\) may be far greater than any of those listed above. Determining neutral densities at large geocentric distances is difficult because interplanetary \(\mbox{Ly}\alpha\) glow intensities exceed those for geocoronal emissions beyond \(8\mbox{--}10~\mbox{R}_{E}\) (Bailey and Gruntman 2011).

Fig. 22
figure 22

Comparison of radial H density profiles from recent analytic models (adapted from Zoennchen et al. 2013)

Recent TWINS observations suggest that the exospheric density may also vary as a function of geomagnetic activity. Bailey and Gruntman (2011) reported \(\sim10\mbox{--}20\%\) increases in the neutral H density at the onset of geomagnetic storms, with the magnitude of the enhancement scaling to the strength of the geomagnetic storm as measured by the minimum in the Dst index. Since the increases in neutral density last for periods on the order of a day or less, the likely source is additional particles on ballistic trajectories with lifetimes on the order of \(13\mbox{--}18~\mbox{hours}\).

3.5 Charge Exchange Within the Magnetosphere

Ring current, Van Allen radiation belt, and plasmaspheric plasmas lie deep within the magnetosphere and exosphere, and are therefore also subject to charge exchange. Although plumes of cold (\(1\mbox{--}10~\mbox{eV}\)) plasmaspheric plasma can extend outward to the magnetopause with densities on the order of \(10~\mbox{cm}^{-3}\) from locations deeper within the magnetosphere where densities are on the order of \(1000~\mbox{cm}^{-3}\), the plasmasphere often terminates abruptly at a sharp plasmapause that lies some \(3\mbox{--}5~\mbox{R}_{E}\) from Earth (Carpenter and Anderson 1992). Some plasmaspheric plasma flows upward through the cusps, joining the population of reflected solar wind ions entering the mantle and magnetotail. A ring current of hot (1–400 keV) plasma with densities ranging from \(1\mbox{--}10~\mbox{cm}^{-3}\) encircles the Earth at radial distances from \(2\mbox{--}7~\mbox{R}_{E}\). Fluxes of energetic (\(>50~\mbox{MeV}\)) ions peak in the inner Van Allen Radiation belt at distances of \(1.2\mbox{--}3~\mbox{R}_{E}\) from the center of the Earth, while fluxes of \(0.1\mbox{--}10~\mbox{MeV}\) energetic electrons peak in the outer Van Allen Radiation belts at radial distances of \(4\mbox{--}5~\mbox{R}_{E}\) from Earth. Radiation belt ion densities are on the order of \(1~\mbox{cm}^{-3}\) (Baumjohann and Treumann 1996; Hultqvist et al. 1999).

Since only high charge state heavy ion populations emit soft X-rays when they exchange electrons with neutrals, we must now consider the composition of the plasmasphere, ring current, and radiation belts to determine whether they are significant sources of soft X-rays. The plasmasphere is comprised of singly charged protons (\(93\mbox{--}97\%\)), He (\(2\mbox{--}6\%\)), and O (1%) ions (Moldwin 1997). The ring current is comprised of protons and, particularly during disturbed geomagnetic storm intervals, singly-charged O ions. Theory predicts that singly-charged ions dominate the radiation belt population, even for solar wind source species with purely high charge state populations (Spjeldvik and Fritz 1978). Although high charge state C and O ions may predominate at high energies (Cohen et al. 2017), particularly during injections (D. Mitchell, personal communication, 2017; Sibeck et al. 1988; Allen et al. 2017) the densities of these ions in the outer magnetosphere are insignificant, on the order of only \(10^{-6}\mbox{--}10^{-3}~\mbox{cm}^{-3}\) (Christon et al. 1994; Allen et al. 2016a,b, 2017). We conclude that magnetospheric particle populations are not a significant source of soft X-rays.

4 Other Sources of Soft X-Rays

The soft X-ray emissions generated by charge exchange with solar wind ions at Venus, Mars, Earth, the Moon, and comets must be distinguished from those generated by charge exchange with solar wind ions in the interplanetary medium, those generated by charge exchange with ions not of solar wind origin at Jupiter, those which result from a host of processes other than charge exchange at the planets (e.g., Bhardwaj et al. 2007), and the cosmic soft X-ray background. This section reviews the other sources that may lie in the line of sight of a soft X-ray telescope. With the aid of this information, we will proceed in Sects. 5 and 7 to predict the images that a wide field-of-view soft X-ray telescope would observe for confirmation by observations in Sect. 6.

4.1 Solar Emissions

The Sun is the brightest source of soft X-rays in the heliosphere. Its thermal plasmas generate both continuum and line emissions. Bremsstrahlung (or “braking radiation”) represents a major source of continuum emission for hot (\(\sim10^{6}~\mbox{K}\)) plasmas such as those found in the Sun’s corona. Brehmstrahlung radiation is generated via the acceleration of charged particles colliding with targets such as atomic nuclei. Ion-electron recombination also results in emissions. In equilibrium plasmas, the ionization that results from the predominant electron-ion collisions is balanced by radiative and dielectronic electron-ion recombination. Both types of recombination generate line emissions whose energy depends on ion charge states and therefore on the ambient plasma temperature. For example, ion species such as \(\mbox{O}^{3+}\) are present in \(10^{5}~\mbox{K}\) plasmas whereas \(\mbox{O}^{7+}\) is present in \(10^{6}~\mbox{K}\) plasmas. The line radiation resulting from recombination lies mainly in the EUV for \(10^{5}~\mbox{K}\) gases, but in the soft X-ray part of the spectrum for \(10^{6}~\mbox{K}\) gases.

4.2 Emissions from the Heliosphere

Interstellar neutrals cross the boundaries of the heliosphere, enter the solar system, and exchange charges with solar wind ions that then generate soft X-rays. Due to the motion of the Sun through the local interstellar cloud, interstellar neutral H (\(\sim85\%\) by composition) and He (\(\sim15\%\)) move with an apparent speed of about \(26~\mbox{km}\,\mbox{s}^{-1}\) relative to observers in the solar system reference frame. Helium atoms appear to flow from an ecliptic longitude and latitude (\(\lambda, \beta\)) of (\(255^{\circ}, 5.5^{\circ}\)), whilst H atoms appear to originate from a slightly different direction (\(252^{\circ}, 9^{\circ}\)) (Lallement et al. 2005). The difference results from hydrogen neutrals exchanging charges with shocked protons at the distorted heliospheric interface, thereby forming a secondary neutral H population with the characteristics of the compressed protons.

As they move towards the inner solar system, neutral interstellar H and He atoms experience the Sun’s effects differently. Neutral H atoms move sunward with the relative motion of the Sun and the Local Cloud and are affected by the attractive force of gravity and repulsive force of radiation pressure. Charge exchange with outward moving solar wind ions results in antisunward-moving neutrals and pick-up ions. Together with solar EUV ionization and electron impact ionization, this charge exchange excludes neutral H from a cavity around the Sun whose \(\sim1\mbox{--}2~\mbox{AU}\) size depends on solar activity through the strength of the depletion processes (Quémerais et al. 2006).

The situation is very different for neutral He. Because radiation pressure has little effect on neutral He, these atoms execute Keplerian hyperbolic orbits to form the He focusing cone downstream from the Sun. EUV solar photons ionize the He atoms, but the resulting ionization cavity extends only about 0.5 AU from the Sun. Consequently, the Earth and spacecraft monitoring the solar wind at the L1 libration point pass through the substantially enhanced neutral He densities within the focusing cone once each northern hemisphere winter (e.g., Dalaudier et al. 1984; Gruntman 1994; Gloeckler et al. 2004). Just as in the case of the H ionization cavity, the densities and sizes of the He focusing cone and cavity depend on solar activity (Lallement et al. 2004b). Figure 23 presents the predicted distributions of interstellar H and He atoms within the heliosphere during solar maximum.

Fig. 23
figure 23

The Sun-centered images showing cuts through the ecliptic plane out to a radius of 10 AU, with longitudes of \(0^{\circ}\) and \(90^{\circ}\) labeled, the location of the Earth noted, and the downwind direction marked by DW. Upper Left: The modeled H density distribution with values running from \(0~\mbox{cm}^{-3}\) (black/purple) to \(6.1\times10^{-2}~\mbox{cm}^{-3}\) (red/white). Upper Right: The modeled He density distribution with values running from \(0~\mbox{cm}^{-3}\) to \(5.2\times10^{-2}~\mbox{cm}^{-3}\). The simulated densities are based on models by Dalaudier et al. (1984) and Lallement et al. (2004a) for He and Lallement et al. (1985a,b) for H. Lower Left: The solar wind proton density as modeled by ENLIL with a logarithmic color scale ranging from \(10^{-2}~\mbox{cm}^{-3}\) to \(10^{3}~\mbox{cm}^{3}\). Lower Right: The relative X-ray emissivity, \(\epsilon=(n_{\text{H}} + Fn_{\text{He}})~n_{swp}~V_{rel}\), where the logarithmic scale runs from \(10^{4}\) to \(10^{8}~\mbox{cm}^{-5}\,\mbox{s}^{-1}\) (Kuntz et al. 2015). \(F\) is a scale factor near unity that accounts for the small difference between the interaction cross sections of H and He

The spatial and temporal variations of the high charge state population within the solar wind introduce further structure into soft X-ray emissions within the heliosphere. As discussed in Sect. 3.3, the largest and strongest heliospheric plasma density variations are those associated with corotating interaction regions and coronal mass ejections. Spiral shock fronts associated with corotating interaction regions provide factor of \(\sim3\) density enhancements that last \(\sim1\) day (Borovsky and Denton 2010), while the factor of 2 to 3 magnetosheath density enhancements (Guo et al. 2010) that precede CMEs last \(\sim11\) hours (Zhang et al. 2008).

Since solar wind charge exchange emissions are proportional to neutral population densities, the expected and observed azimuthal asymmetry in the heliospheric neutral He densities implies an asymmetry in heliospheric soft X-ray emission, a topic further explored in Sect. 5.4. The arrival of solar wind structures with different densities and compositions results in time- and spatial dependent variations in soft X-ray emissions superimposed upon those due to the neutral density asymmetries. Since the emission spectrum is comprised of many lines from 150 to 350 eV (e.g., see Fig. 15), it is difficult for detectors with low or medium spectral resolution to distinguish this spectrum from the thermal spectrum of the local bubble (see Sect. 4.4).

4.3 Planetary Emissions

Processes other than solar wind charge exchange can generate soft X-rays at Venus, Mars, Earth, and Earth’s Moon. Jupiter and Saturn present particularly interesting cases. Here charge exchange produces soft X-ray aurorae, electron bremsstrahlung dominates the auroral spectra at energies above 3 keV, and the brightness of the planetary disk in soft X-rays varies proportionally to that of solar X-rays. The latter is particularly noticeable when individual solar flares are mirrored in the Jovian soft X-ray light curve.

The Einstein Observatory provided the first detection of X-rays with energies \(0.2\mbox{--}3.0~\mbox{keV}\) from Jupiter’s aurora. Metzger et al. (1983) proposed that this emission is related to energetic ion precipitation and noted that either a combination of O and S line emissions or electron bremsstrahlung continuum could fit the spectral data. Since the electrons probably cannot input sufficient power and the spectra are too soft for X-ray emission due to bremsstrahlung, line emissions from O and S heavy ion precipitation was taken to be a more likely cause of the X-ray aurora. ROSAT observations in the early 1990s confirmed this general picture (Waite et al. 1994). Horanyi et al. (1988) initially modeled the auroral X-ray emissions in terms of precipitating low (\(q<4\)) charge state O and S ions. Cravens et al. (1995) subsequently invoked charge exchange with highly charged O ions.

The precipitating ions could originate in the magnetosphere (e.g., Io’s volcanoes) or solar wind. It should be possible to distinguish between these sources by inspecting the charge exchange emission lines in the \(0.3\mbox{--}0.4~\mbox{keV}\) band, where the presence of S lines would indicate an Io origin, whereas C lines would indicate a solar wind origin. Cravens et al. (2003) concluded that both cases require substantial particle acceleration to produce the observed X-ray fluxes.

Surprisingly, Chandra X-ray observatory (Weisskopf et al. 2002) observations of Jupiter in 2000 revealed that the polar hot spot of X-ray emissions pulsates with a well defined 45 minute period and maps magnetically to distances exceeding \(30~\mbox{R}_{\mathrm{J}}\) from the planet (Gladstone et al. 2002). Subsequent Chandra observations indicate much less organized periodicities ranging from 20 to 70 min (Elsner et al. 2005). Bunce et al. (2004) attributed these periodicities to particle acceleration driven by pulsed reconnection at the dayside magnetopause.

For “fast flow” solar wind conditions with high plasma densities and IMF strengths, and potential drops of \(\sim5~\mbox{MV}\) that strip the ions (e.g., Cravens et al. 2003), precipitating magnetospheric O ions can produce X-ray intensities that match the observations. The heavy ion precipitation should be associated with fluxes of relativistic electrons escaping Jupiter (Ozak et al. 2013) as well as significant downward field-aligned currents. The current Juno mission will no doubt shed much light on magnetosphere-ionosphere interactions at Jupiter.

XMM-Newton (Jansen et al. 2001) and Chandra X-ray observations demonstrate that line emissions (in particular \(\mbox{O}^{6+}\)) resulting from charge exchange between highly stripped energetic ions and \(\mbox{H}_{2}\) molecules in the planet’s upper atmosphere dominate Jupiter’s auroral X-ray spectrum below 2 keV (Branduardi-Raymont et al. 2004, 2007; Elsner et al. 2005). At higher energies (\(2\mbox{--}10~\mbox{keV}\)), the featureless auroral X-ray spectrum can be attributed to electron bremsstrahlung. Figure 24 shows XMM-Newton EPIC spectral maps in narrow energy bands: the aurorae are very evident in the lower energy band centered on the charge exchange \(\mbox{O}^{7+}\) line (top left panel) and the higher energy bands where bremsstrahlung dominates (bottom panels). By contrast, a round uniform disk is observed for a band centered on the Fe lines (top right panel) that characterize the scattered solar coronal spectrum. Unfortunately neither XMM-Newton nor Chandra possess the combination of collecting area and high resolving power needed to separate C from S lines in the \(0.3\mbox{--}0.4~\mbox{keV}\) band. There are some indications from XMM-Newton and Chandra (Branduardi-Raymont et al. 2007; Hui et al. 2009, 2010a; Ozak et al. 2010) that S lines may provide a better spectral fit at low energies, implying ions of magnetospheric origin.

Fig. 24
figure 24

XMM-Newton EPIC CCD images of Jupiter in narrow energy bands. The top left and two bottom panels show auroral contributions in the \(0.55\mbox{--}0.60~\mbox{keV}\) \(\mbox{O}^{6+}\) line, at \(3\mbox{--}5~\mbox{keV}\), and at \(5\mbox{--}10~\mbox{keV}\), while the top right panel shows the disk contribution in the \(0.70\mbox{--}0.75~\mbox{keV}\) and \(0.80\mbox{--}0.85~\mbox{keV}\) \(\mbox{Fe}^{17+}\) lines (Branduardi-Raymont et al. 2007). The color scale bar is in units of EPIC counts

Electrons are believed to produce ultraviolet auroral features at Jupiter. As shown in Fig. 25, Chandra observations demonstrate that high (but not low) energy X-ray emissions tend to occur over the ultraviolet auroral oval and other bright ultraviolet features, suggesting that the 10–100 keV precipitating electrons that excite atmospheric H and H2 also produce high energy X-ray emissions (Branduardi-Raymont et al. 2008). Since Jupiter’s ultraviolet aurora brightened following the arrival of a solar wind shock (Clarke et al. 2009), it would not be too surprising for the electron bremsstrahlung X-ray emission to follow the same trend. And indeed, the electron bremsstrahlung spectral component varied significantly over a 3.5 day XMM-Newton observation in 2003 November following the “Halloween Storm”, probably in response to magnetospheric particle energization caused by a compression of the Jovian magnetosphere corresponding to the arrival of a solar wind shock or a CME.

Fig. 25
figure 25

Polar projection of Chandra soft charge exchange (\(<2~\mbox{keV}\), small green dots) and hard electron bremsstrahlung (\(>2~\mbox{keV}\), large green dots) X-ray emission events superimposed on a simultaneous Hubble Space Telescope Imaging Spectrograph UV image (orange) of Jupiter’s aurora (Branduardi-Raymont et al. 2008). The \(10^{\circ}\) spaced grid is fixed in System III with longitude \(180^{\circ}\) toward the bottom and longitude \(90^{\circ}\) to the right. The system III \(z\)-axis lies along the planet’s spin axis and longitude increases from East to West according to an observer at Earth

Chandra observed a brightening of auroral X-ray emissions, mostly below 0.5 keV, around the time when a CME was expected to impact Jupiter in October 2011. As indicated by the map shown in Fig. 26, a comparison with magnetic field models (Vogt et al. 2011) indicates that these emissions tend to cluster at the footprints of open magnetic field lines that map to the outer magnetosphere. This led Dunn et al. (2016) to associate the origin of at least some of the X-ray emissions with possible direct solar wind O and C precipitation.

Fig. 26
figure 26

Chandra X-ray events (colored dots, plotted in Jupiter’s System III coordinate system: colors and species are shown on the RHS) gather at the footprints of open field lines in Vogt et al. (2011) model (colored oval). The red arrow indicates the Sun/noon position (from Dunn et al. 2016)

Saturn, like Jupiter, emits powerful auroral emissions at radio, infrared, and ultraviolet wavelengths. Solar wind compressions cause ultraviolet and radio brightenings (e.g., Clarke et al. 2009). By analogy with Jupiter, X-ray aurorae powered by charge exchange should also be expected on Saturn, yet none have been observed to date. No auroral X-rays were observed at the time when a CME was predicted to arrive (Branduardi-Raymont et al. 2013), perhaps because the required accelerating potentials were absent (Hui et al. 2010b). Nevertheless, the planet is a source of soft X-rays. A combination of elastic and fluorescent scattering of solar X-rays in the H2 and CH4 atmosphere can account for Saturn’s disk, polar cap and ring emissions.

As for Jupiter, the flux from the disk tracks that of solar X-rays (Bhardwaj et al. 2005a,b; Branduardi-Raymont et al. 2010). Fluorescence also explains most of the disk emissions observed from Venus and Mars. Here solar X-rays ionize and remove K-shell electrons from the C and N atoms in atmospheric neutrals like CO2 or N2 (while on Jupiter and Saturn this occurs for the C in CH4). The emissions occur when the K-shell vacancy is filled by a higher energy valence electron.

Looking further afield, we can expect Uranus, Neptune, and Saturn’s moon Titan to emit X-rays by magnetospheric particle precipitation, and by scattering of solar X-rays. Nevertheless, Bhardwaj et al. (2007) reported that Chandra failed to detect X-rays from Uranus in August 2002. A comparison of the parameters relevant to aurora production at Uranus and Neptune with those for Jupiter leads to the conclusion that X-ray emissions from these planets are far too faint to be detected by current X-ray observatories, but may be just bright enough to be observed by Athena, the next generation X-ray observatory, when it flies around 2028 (Branduardi-Raymont et al. 2010). There is even some evidence for X-ray emission from Pluto (Lisse et al. 2017).

In addition to charge exchange, there are other non-thermal mechanisms for generating soft X-rays in solar system environments much colder than the solar corona. These environments include the Earth’s upper atmosphere, comets, and the Jovian upper atmosphere where neutral temperatures are only \(\sim 1000\), 20, and 300 K, respectively. Electron temperatures in the ionospheres of these bodies are somewhat higher than these neutral temperatures, but not by more than a few thousand degrees. Consequently, the X-ray emissions from these targets result from non-thermal processes and not thermal collisions. For example, the precipitation of electrons accelerated to high energies in the Earth’s magnetotail produces bremsstrahlung X-rays in the Earth’s auroral oval. In addition, planetary atmospheres scatter solar X-rays through both Thompson and fluorescent processes (e.g., Schmitt et al. 1987; Snowden and Freyberg 1993).

4.4 Cosmic Soft X-Ray Emissions

The cosmic X-ray sky comprises numerous components that are strongly spectrally and spatially variable. The emissions from some of these components are comparable to or brighter than those typically generated by the solar wind charge exchange processes that occur much closer to Earth. To conduct the science outlined within this paper, these background emissions must be quantified and then subtracted from soft X-ray images. Here, we consider diffuse, point, and distinct extended sources (see Fig. 27).

Fig. 27
figure 27

False color image of the soft X-ray background with Red: \({1\over 4}~\mbox{keV}\) emission, Green \({3\over 4}~\mbox{keV}\) emission, and Blue 1.5 keV emission. The plot is in Galactic coordinates centered on the Galactic Center and the coordinate grid marked every 30 degrees. Stars indicate the north ecliptic pole (in the northern Galactic hemisphere) and the south ecliptic pole (in the southern Galactic hemisphere). The bulk of the X-ray point sources have been removed; the remaining point-like sources near the Galactic plane are mostly supernova remnants, while those at high latitudes are usually clusters of galaxies. Outlined or otherwise identified are the Virgo cluster of galaxies, the Loop I superbubble, the Eridanus-Orion bubble (otherwise known as the Eridion bubble), the Cygnus superbubble, the Crab, Vela, and Cygnus Loop supernova remnants, the Galactic X-ray bulge, the nearby galaxy LMC, and the Sco X-1 neutron star (the brightest X-ray source in the sky as seen from Earth). The Galactic halo emission is the red emission at Galactic latitudes above \(30^{\circ}\)

First observed in the late 1960s (e.g, Bowyer et al. 1968), non-heliospheric diffuse emissions originate at locations ranging from the local interstellar medium to cosmological distances (e.g., McCammon and Sanders 1990). Galactic emissions from a thermal plasma within the Local Hot Bubble (e.g., Tanaka and Bleeker 1977; Sanders et al. 1977; Snowden et al. 1990, 2014; Snowden 2002), a low density (\(0.05~\mbox{cm}^{-3}\)), high temperature (\(10^{6}~\mbox{K}\)), region extending from \(\sim30\) to \(\sim150\) parsecs from the Sun depending on direction, dominate the flux of soft X-rays at lower (\(\sim{1\over 4}~\mbox{keV}\)) energies. High neutral H column densities within the Milky Way disk absorb emissions from greater distances. There are contributions from the lower Galactic halo at high Galactic latitudes (e.g., Kuntz and Snowden 2000). Consequently, as shown in the upper panel of Fig. 28 (Snowden et al. 1997), the surface brightness at \({1\over 4}~\mbox{keV}\) generally increases from the Galactic equatorial plane towards both the north and south poles. While some of the features superimposed upon this general pattern can be identified with specific Galactic objects, most result from integral LOS filling factor and density (i.e., emission-measure) variations in diffuse \(\sim10^{6}~\mbox{K}\) plasmas further strongly modified by variable absorption in the interstellar medium.

Fig. 28
figure 28

ROSAT All-Sky Survey (RASS, Snowden et al. 1997) images of the soft X-ray background shown as Aitoff-Hammer equal-area maps in Galactic coordinates centered on the Galactic center. Galactic longitude increases to the left, the south Galactic pole is at the bottom and the north Galactic pole is at the top. Purple and blue indicate low intensity while red and white indicate high intensity. The units of the color bars are ROSAT \(\mbox{counts}\,\mbox{s}^{-1}\,\mbox{arcmin}^{-2}\). Upper Panel: \({1\over 4}~\mbox{keV}\) band; Middle Panel: \({3\over 4}~\mbox{keV}\) band; Lower Panel: 1.5 keV band. The white circle in the lower right of all figures outlines a \(10^{\circ}\) radius region surrounding the south ecliptic pole

At the higher energies (\({3\over 4}~\mbox{keV}\)) shown in the middle panel of Fig. 28, distinct objects (some extending over large solid angles) dominate a relatively flat background. The strong enhancement at the Galactic center represents the combined emission from the nearby Loop I Superbubble and the Galactic Bulge (e.g., Snowden et al. 1997). The generally flat background combines the extragalactic background (primarily unresolved point sources) with emissions from the Galactic halo and our local group of galaxies modulated by absorption from the interstellar medium near the Galactic plane. The 1.5 keV band map is shown in the lower panel of Fig. 28. It shows a structure similar to the \({3\over 4}~\mbox{keV}\) band map except that the distinct features are not as bright.

As can be seen in the maps, the cosmic diffuse X-ray background (e.g., Snowden et al. 1997) is bright and spatially varying, differs radically in the \({1\over 4}~\mbox{keV}\) and \({3\over 4}~\mbox{keV}\) bands, and is observed in all directions. It will therefore be present in all observations of soft X-rays generated by solar wind charge exchange. However, it is temporally constant on all human time scales and is well understood and mapped. Consequently it can be subtracted from the light curves and images of soft X-rays generated by charge exchange observations. While the surface brightness of the cosmic background varies by up to an order of magnitude when small regions are considered, the south ecliptic pole (a likely background direction for a soft X-ray mission imaging the subsolar bow shock and magnetopause from a polar vantage point) lies in a relatively benign direction, particularly at \({1\over 4}~\mbox{keV}\). At \({3\over 4}~\mbox{keV}\), the Large Magellanic Cloud (a nearby galaxy seen in the right panel of Fig. 29) does show enhanced emissions but it has been particularly well studied in soft X-rays (e.g., Snowden and Petre 1994; Haberl 2014).

Fig. 29
figure 29

Left Panel: XMM-Newton mosaic of the Coma Cluster of galaxies in the \(0.4\mbox{--}1.25~\mbox{keV}\) band. The data are square root scaled and the units are in \(\mbox{counts}\,\mbox{s}^{-1}\,\mbox{deg}^{-2}\). The coordinates are in right ascension and declination. Right Panel: XMM-Newton mosaic of a region of the Large Magellanic Cloud, a satellite galaxy of the Milky Way, in the same band, coordinates, and units (unpublished images provided by S.L. Snowden)

In addition to the diffuse emission, there are many point sources, including stars, compact objects (e.g., pulsars, X-ray binaries), and active galactic nuclei (AGN). Figure 30 shows the locations of the 18,811 bright sources (Bright Source Catalog, BSC) detected during the ROSAT All-Sky Survey (RASS) (Voges et al. 1999). A further 105,924 sources were included in the Faint Source Catalog extension to the BSC. Much of what appears to be a general diffuse emission at higher energies in Fig. 28 (e.g., at high latitudes in the middle and lower panels) actually results from the superposition of unresolved emission from AGN at cosmological distances. These individual sources are insufficiently bright to be detected with the available exposure time allowed by the RASS, but can be resolved by newer observatories particularly during deep observations like the XMM-Newton image of the Hubble Deep Field North shown in the right panel of Fig. 31. As shown in the upper panel of Fig. 28, distinct galactic supernova remnants (SNRs) such as Vela and Puppis (enlarged in the left panel of Fig. 31) can subtend relatively large areas on the sky and can be both bright and strongly spatially varying. Nearby galaxies and clusters of galaxies can contribute as both point and extended sources, sometimes spectacularly as illustrated by the examples in Fig. 29.

Fig. 30
figure 30

Locations, relative fluxes (the size of the circle scales as the log of the flux), and hardnesses (purple indicates a hard source, red a soft source) of the 18,811 sources listed in the RASS Bright Source Catalog (Voges et al.