Interplanetary coronal mass ejections (ICMEs) are large-scale heliospheric transients that originate from the Sun. When an ICME is sufficiently faster than the preceding solar wind, a shock wave develops ahead of the ICME. The turbulent region between the shock and the ICME is called the sheath region. ICMEs and their sheaths and shocks are all interesting structures from the fundamental plasma physics viewpoint. They are also key drivers of space weather disturbances in the heliosphere and planetary environments. ICME-driven shock waves can accelerate charged particles to high energies. Sheaths and ICMEs drive practically all intense geospace storms at the Earth, and they can also affect dramatically the planetary radiation environments and atmospheres. This review focuses on the current understanding of observational signatures and properties of ICMEs and the associated sheath regions based on five decades of studies. In addition, we discuss modelling of ICMEs and many fundamental outstanding questions on their origin, evolution and effects, largely due to the limitations of single spacecraft observations of these macro-scale structures. We also present current understanding of space weather consequences of these large-scale solar wind structures, including effects at the other Solar System planets and exoplanets. We specially emphasize the different origin, properties and consequences of the sheaths and ICMEs.
Coronal mass ejections; Interplanetary shocks; Magnetic clouds; Solar wind; Space weather.
© The Author(s) 2017.
(Left) A schematic of an ICME showing also the leading fast forward shock (arc), and the sheath region. The ICME shown here is depicted to have a flux rope structure, not always detected in-situ. (Right) Solar wind observations during an ICME from the ACE spacecraft located at the Lagrangian point L1. The panels show from top to bottom: the magnetic field magnitude, the longitude and latitude angles of the magnetic field in the Geocentric Solar Magnetospheric (GSM) coordinate system, and the solar wind speed. The blue dashed line marks the shock and the ICME is bounded by the pair of red lines
(Left) An ICME with a magnetic cloud (MC) structure and (right) a non-magnetic cloud ICME. The panels show from top to bottom:
a magnetic field magnitude, the b longitude and c latitude angles of the magnetic field in GSM coordinates, ( ϕ B = 90 ∘ ) is defined eastward (westward), and ϕ B = 270 ∘ ( θ B = + 90 ∘ ) is defined northward (southward), solar wind θ B = - 90 ∘ d speed, e proton temperature (black: measured temperature, red: expected temperature, see Sect. 2.1), f plasma beta, g helium to proton ratio, h O O + 7 / ratio (black) and Fe/O ratio (red), + 6 i iron average charge state, and j pitch angle spectrogram of 272-eV electrons. Pitch angle ( 0 ∘ ) refers to the particles that stream parallel (anti-parallel) to the magnetic field. A pair of red solid lines bound the ICME interval and the dashed blue line shows a shock. The red dashed lines bound the interval with MC signatures. The times are from the Richardson and Cane ICME list. The measurements are from the ACE spacecraft. Data have been obtained through the ACE Data Center ( 180 ∘ http://www.srl.caltech.edu/ACE/ASC/)
Forbush decrease associated with a shock-driving ICME. Image reproduced by permission from Cane (2000), copyright by Kluwer
Different spacecraft paths (Group 1–Group 3) through an ICME and the corresponding perpendicular pressure (
) profiles. The different groups are described in the text. The contours in the figure show the density and the numbers give the density ratios between the solar wind and magnetopause values. In the right-hand part of the figure the first dashed line shows the shock and the following two dashed lines bound the ICME interval. Images reproduced by permission from Jian et al. (2006), copyright by Springer P t
Flux rope categories for Top) bipolar (low-inclination) and Bottom) unipolar (highly-inclined) magnetic clouds. The letters above each flux rope defines the direction of the magnetic field at the leading edge of the MC, at its centre, and at its trailing edge (E
East, W = West, N = North, S = South), assuming that the spacecraft moves into the page. The sign of magnetic helicity is either left-handed/negative (LH; clockwise rotation) or right-handed/positive (RH; counter-clockwise rotation). Image courtesy of Erika Palmerio =
Examples of sheath–ICME boundaries:
a magnetic field magnitude, b components of the magnetic field in GSE coordinates ( : purple, B X : green, B Y : red), solar wind B Z c speed, d proton temperature, and e density. In the left-hand panels the dashed vertical line indicates the sharp boundary and in the right-hand panels the pair of dashed lines marks a wider boundary layer. Both cases are discussed in Wei et al. (2003a)
Different regions of the ICME. The panels show the magnetic field
a magnitude, b latitude, and c longitude, d solar wind speed, e density f temperature (black: measured, red: expected temperature), g total pressure perpendicular to the magnetic field, h plasma beta, i pitch angle spectrogram of 272-eV electrons, j red: helium to proton ratio, black: O /O + 7 ratio, and + 6 k average iron charge state. The data are from the ACE spacecraft. Images reproduced by permission from Kilpua et al. (2013b), copyright by the authors
(Left) The global flux rope axis configuration as deduced from multi-spacecraft observations for a magnetic cloud observed on January 6–8, 1978. All these spacecraft were located close to the ecliptic plane at radial distances between 1–2 AU. (Right) Schematic of a flux rope loop that is distorted along the Parker spiral and carries a sector boundary crossing. Images reproduced by permission from [left] Burlaga et al. (1990); [right] from Crooker et al. (1998), copyright by AGU
Snapshot of the CME evolution for an MHD simulation at four different times. The colour map shows the radial velocity, the black contours the magnetic flux, and the red contours the number density. Image reproduced by permission from Riley and Crooker (2004), copyright by AAS
Example event from the HELCATS database. (Left) STEREO-B COR2 coronagraph image of an Earth-directed CME that erupted on July 9, 2013. (Middle) J-map that has been constructed using STEREO-B heliospheric imager observations. Red dots show the model fitted values and they indicate the path of the CME. (Right) ACE in-situ observations of the associated ICME. The panels are the same as in Fig. 1. The blue dashed line marks the shock and the ICME is bounded between the pair of red lines
(Left) A CME with a classical three-part structure detected one February 27, 2000 in SOHO/LASCO C3 image (Image courtesy: NASA). (Right) An excess mass image illustrating two more signatures of a five-part CME on June 10, 2000 in SOHO/LASCO C2 image. The yellow arrow indicates the sharp edge (shock) and the green arrow the bright front. Image reproduced by permission from Vourlidas et al. (2013), copyright by Springer
Fitting of the force-free field cylindrical symmetric constant-
model by Burlaga (1988) to the observed magnetic field data using the least-squares fitting method developed by Lepping et al. (1990). The panels show (top) magnetic field magnitude, and the magnetic field (middle) latitude and (bottom) longitude in solar ecliptic coordinates. α and ϕ 0 denote the flux rope axis orientation from the fitting, and θ 0 is the radial diameter of the flux rope. 2 R 0 is the parameter that is related to the goodness of the least-squares fit of the magnetic field data in the model (the smaller the value, the better the fit). Image reproduced by permission from Lepping et al. (1990), copyright by AGU ξ 2
(Left) Results of the fitting of a magnetic cloud on February 7, 1981 using the Lundquist (“circular”) and elliptical (“oblate”) models (the smooth black lines). Both models include the effect of the expansion. (Right) Fitting results for two magnetic clouds from a non-force-free model that incorporates also the hydrostatic plasma pressure and the proton current density in the fitting procedure (the smooth red lines). Images reproduced by permission from [left] Vandas et al. (2006), copyright by COSPAR; [right] Hidalgo (2016), copyright by AAS
Examples of Grad–Shafranov reconstruction of two magnetic clouds observed by the Wind spacecraft. (Top) Flux rope cross-sectional shapes in the plane perpendicular to the invariant axis. The black arrows show the spacecraft magnetic field observations projected on the same place. The white thick countour corresponds to the flux rope boundary and the white dot the axis of the rope. The colour coding gives the magnitude of the magnetic field component along the invariant axis direction (
z). (Bottom) curves. The red (green) curves represent P t ( A ) A along the inbound (outbound) parts of the spacecraft trajectory through the MC. The results are from the European Union’s FP7 project HELCATS online catalogues ( http://helcats-fp7.eu)
Geometries for a magnetic cloud of 19 March 2001 as obtained from
a torus model and b cylindrical model. The arrows denote the direction of magnetic field on the surface of the (S), the direction of the axial field, and (A) and the spacecraft trajectory (S/C). c Curved flux rope loop featuring the central but (path A) and the flank encounter (path F) with a double crossing through the axis. d Spheroidal oblate flux rope. The solid (dashed) lines depict the magnetic field lines above (below) the plane of the figure. The thick line shows the reference ellipsoid. Panels ( a)–( c) reproduced by permission from Marubashi and Lepping (2007), copyright by the authors; panel ( d) from Vandas et al. (1993), copyright by AGU
(Left) Examples of field line turns in a magnetic cloud of August 30, 2004 (Event 7 in the right-hand plot). (Right) The average field line twist versus the shifted flux function
, where | A - A 0 | is the value at the flux rope centre and A 0 A is the GSR result. The square indicates the mean for each of the curves, and the error bars give the standard deviations. Image reproduced by permission from Hu et al. (2015), copyright by AGU
An example of a sheath region between a forward shock and an MC. The panels show from top to bottom:
a the magnetic field magnitude, b components of the magnetic field in GSE coordinates ( : purple, B X : green, B Y : red), solar wind B Z c speed, d proton temperature, and e proton density. The dashed line shows the shock and the solid line the leading edge of the ICME. The measurements are from the ACE spacecraft. Data have been obtained from the ACE Data Center ( http://www.srl.caltech.edu/ACE/ASC/)
Probability distributions of various solar wind parameters in ICME sheaths (black) and ICME (red). The panels give:
a magnetic field magnitude, b IMF north-south component, c root-mean-square of the magnetic field, solar wind d speed, e density, f temperature, g dynamic pressure, h plasma beta, i Alfvén Mach number, and j alpha to proton (He /p), + + k O / + 7 , and + 6 l Fe/O ratios. Panels a– i The data sets are 5-min OMNI data, panel j 1-h OMNI data and panels k– l are 1-h (2-h after August 2011) ACE/SWICS data. The numbers in parenthesis show the number of sheaths and ICMEs
ULF power of the north-south component of IMF in ICME sheaths. The black line is the median. The red and blue lines are upper and lower quartiles, respectively. The length of all sheaths is scaled to 10 h (see Kilpua et al. 2013a)
Draping of the IMF around the ICME. Sketch of the draping of the IMF around and ICME in
a the ecliptic plane and b the meridional plane. c Global MHD simulation results of a Parker spiral type IMF draping. The strength of the out-of-ecliptic IMF component is shown for a view towards the Sun (ignore the part within the oval that belongs to ICME). Panels ( a) and ( b) reproduced by permission from Gosling and McComas (1987); Panel ( c) from Siscoe et al. (2007), copyright by AGU
Particle spectra of a gradual SEP event (left) and two consecutive impulsive events (right) as observed by the ISEE 3 spacecraft. Image reproduced by permission from Reames (1999), copyright by Springer
Annual variations in the ICME (black) and sheath (blue) occurrence rate and properties. The panels show
a the yearly mean sunspot number from Solar Influences Data Center ( http://sidc.oma.be), and the annual number of b ICMEs, and c the ICMEs that drove shocks, the annual means of d the duration, and the peak e magnetic field, and f speed in ICMEs (black) and sheaths (blue). The error bars give the standard deviations. The years when three or less events occurred are excluded from panels ( d– f). The ICME intervals are from the Richardson and Cane ICME list. The sheath intervals are determined using the ICME leading edge times given in Richardson and Cane list and the shocks time from the Interplanetary Shock Database of the University of Helsinki ( www.ipshocks.fi/). Note that here we consider as “sheaths” only the cases where an ICME was preceded by a fully developed shock
The top panel shows the monthly sunspot number. The next three panels show the percentage (three solar rotation averages) of the solar wind flows associated with ICMEs, high-speed streams and solar solar wind for over more than four solar cycles (1963–2011). The figure is from Richardson and Cane (2012a)
Solar cycle variations of NS- and SN-type bipolar magnetic clouds, and bipolar (S and N types combined) and unipolar (SN and NS types combined) magnetic clouds in the near-Earth solar wind. The four first panels show the annual counts and the fifth panel shows the N and S tilt angle of the heliospheric current sheet. The flux rope types are illustrated in Fig. 5. In this figure the letters refer only to the north-south magnetic field component (e.g., N can present both ENW and WNE type flux ropes). Image reproduced by permission from Li et al. (2011), copyright by Springer, where the third panel was erroneously labeled as bipolar, while it clearly shows the northward unipolar clouds
An ICME detected by Wind at 1 AU, Ulysses at 5.3 AU and Voyager 2 at 58 AU (bottom panel) using the total pressure (magnetic + plasma) perpendicular to the magnetic field (not available from Voyager 2) and alpha to proton ratio (He
H + + / ). The data from different spacecraft are time-shifted to align the ICMEs. Image reproduced by permission from Richardson et al. (2006); copyright by COSPAR (see also references therein) +
The width of ICMEs as a function of the heliospheric distance
R from the Sun. Diamonds show the radial widths of ICMEs that were combined by Richardson et al. (2006) using previously published ICME lists. The widths averaged over 3 AU are shown by horizontal bars and the vertical bars give the error of the mean. The dashed line gives the best linear fit to observations inside 15 AU where ICMEs generally still expand. Image reproduced by permission from Richardson et al. (2006), copyright by COSPAR
The panels show from top to bottom the average solar wind density, speed, temperature and magnetic field magnitude as a function of heliospheric distance (
R) from the Sun. Also shown are fits to the (solid) ICME data and fits to the (dashed) ambient, i.e., non-ICME, solar wind data. The data points combine measurements from Helios 1 and 2, Wind, ACE, Ulysses and Voyager spacecraft. Image reproduced by permission from Richardson et al. (2006), copyright by COSPAR; see also the references therein
(Left) Voyager 2 magnetic field data measured within the ICME sheath region at the heliospheric distance of 73.2 AU from the Sun. The panels give from top to bottom the magnetic field components in the RTN coordinates and the magnetic field magnitude. Intervals A and B are planar structures, while the horizontal blue bar indicates the interval where PMSs were not found. (Right) The
– ϕ diagram (where θ and ϕ are the IMF longitude and latitude). The plane perpendicular to the PMS normal is represented by the dashed curve. PMSs are identified as the periods when magnetic field vectors are distributed in the close proximity of the curve presented. The top panel shows the interval A when the PMSs were found, and the bottom features the non-PMS interval. Images adapted from Intriligator et al. (2008), copyright by AGU θ
A complex ejecta. The panels show from top to bottom: solar wind speed, temperature, density, alpha to proton and O
O + 8 / ratio, elevation and azimuth angle of the magnetic field, and magnetic field magnitude. Image reproduced by permission from Burlaga et al. (2002), copyright by AGU + 7
Illustration of the magnetic configuration of (left) a non-eroded magnetic cloud and (right) an eroded magnetic cloud. The red curve shows the variations of the accumulated azimuthal flux and the blue and black lines the variations in the field components. The coordinate system is the MC frame (see definition, e.g., in Dasso et al. 2005). Image reproduced by permission from Ruffenach et al. (2012), copyright by AGU
Example of a sheath and a magnetic cloud driving geomagnetic activity on December 14–16, 2006. The panels show from top to bottom:
a magnetic field magnitude, b magnetic field north-south component (in GSM coordinates), c solar wind speed, d solar wind dawn-dusk electric field, e Dst, f AE (see Sect. 6.2) and g Kp. The data are 1-min OMNI data, Dst is 1-h average and Kp 3-h average. The dashed red line marks the shock and the pair of solid lines marks the MC interval
Drivers of geomagnetic storms around solar minimum and solar maximum accumulated over four solar cycles (1964–2011). The geomagnetic storm definition follows NOAA storm scale (
http://www.swpc.noaa.gov/noaa-scales-explanation). Image reproduced by permission from Richardson and Cane (2012b), copyright by the authors
(Left) Drivers of magnetic storms for three different Dst limits. (Right) The minimum Dst versus the maximum Kp for intense storms (Kp
7 and/or Dst ≥ nT) for 1997–2003. Note that in several cases the storm is defined as “intense” only by one of the used indices, while the other indicates more moderate activity (see discussion by Huttunen et al. and Huttunen and Koskinen 2004). Sheath storms are shown by asterisks and magnetic cloud storms by crosses. The rectangles show the cases that did not fulfill the Gonzalez et al. (1998) intense Kp criterion for an intense storm, i.e., the requirement that Kp has to be < - 100 at least for three 3-h periods. Images reproduced by permission from [left] Huttunen and Koskinen (2004), copyright by the authors; [right] Koskinen and Huttunen (2006), copyright by Springer ≥ 6
High-latitude response to sheaths (black diamonds), MCs (red stars), and the MC front (green triangles) and rear regions (blue squares). The data set consist of the same 79 events from 1996 to 2009 as analyzed in Kilpua et al. (2013b). The plot shows the maximum AE as a function of total time when AE
nT within each region. The large symbols show the averages for each region. The first number gives the total number of each region (note that not all MCs have sheaths, front or rear regions) and the second number the cases that had significant high-latitude response (AE > 1000 nT) (percentage also shown in parenthesis) > 1000
Combined long-term observations as measured by SAMPEX (in low Earth polar orbit) and Van Allen Probes (in the equatorial plane; the shaded period from 2012 on). Two top panels show the 27-day window-averaged
-MeV electron fluxes from geostationary spacecraft GOES, and the minimum value of the 1-h Dst index each month. The third panel shows the yearly window-averaged sunspot numbers (black curve) and the weekly window-averaged solar wind speed (km s > 2 ), red curve). The bottom panels show 27-day averaged - 1 -MeV electron flux (units cm ∼ 2 s - 2 sr - 1 MeV - 1 ) combined from SAMPEX and the REPT instrument on the Van Allen Probes. Image reproduced by permission from Li et al. (2017), copyright by the authors - 1
Atmospheric loss at Mars during encounters with ICMEs (red diamonds, blue show the median) compared to integrated loss through the mission duration (solid: median, dashed: first/fourth quartiles) as a function of solar zenith angle. Image reproduced by permission from Jakosky et al. (2015), copyright by AAAS
Blue-wing enhancements (shown by arrows) in dynamic
Balmer line spectra of a fully-convective M4 dwarf star V374 Pegasus occurred. The darkening colours indicate the increasing intensity of the H α region. The vertical axis shows the time in Heliocentric Julian Days (HJD), starting from 2453603.70. Image reproduced by permission from Vida et al. (2016), copyright by ESO H α
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