Short-term variability of the Sun-Earth system: an overview of progress made during the CAWSES-II period
© Gopalswamy et al. 2015
Received: 16 December 2014
Accepted: 24 April 2015
Published: 17 June 2015
This paper presents an overview of results obtained during the CAWSES-II period on the short-term variability of the Sun and how it affects the near-Earth space environment. CAWSES-II was planned to examine the behavior of the solar-terrestrial system as the solar activity climbed to its maximum phase in solar cycle 24. After a deep minimum following cycle 23, the Sun climbed to a very weak maximum in terms of the sunspot number in cycle 24 (MiniMax24), so many of the results presented here refer to this weak activity in comparison with cycle 23. The short-term variability that has immediate consequence to Earth and geospace manifests as solar eruptions from closed-field regions and high-speed streams from coronal holes. Both electromagnetic (flares) and mass emissions (coronal mass ejections - CMEs) are involved in solar eruptions, while coronal holes result in high-speed streams that collide with slow wind forming the so-called corotating interaction regions (CIRs). Fast CMEs affect Earth via leading shocks accelerating energetic particles and creating large geomagnetic storms. CIRs and their trailing high-speed streams (HSSs), on the other hand, are responsible for recurrent small geomagnetic storms and extended days of auroral zone activity, respectively. The latter leads to the acceleration of relativistic magnetospheric ‘killer’ electrons. One of the major consequences of the weak solar activity is the altered physical state of the heliosphere that has serious implications for the shock-driving and storm-causing properties of CMEs. Finally, a discussion is presented on extreme space weather events prompted by the 23 July 2012 super storm event that occurred on the backside of the Sun. Many of these studies were enabled by the simultaneous availability of remote sensing and in situ observations from multiple vantage points with respect to the Sun-Earth line.
KeywordsSolar activity Space weather Coronal mass ejections Flares Solar energetic particle events Geospace impact Geomagnetic storms
The second phase of the Climate and Weather of the Sun-Earth System (CAWSES-II) was organized into task groups (TGs). Task Group 3 (TG3) was focused on the short-term variability of the Sun-Earth system. Solar variability on time scales up to 11 years was relevant to TG3. The relevant variability occurs in the mass and electromagnetic outputs of the Sun. The mass output has three forms: the solar wind, coronal mass ejections (CMEs), and solar energetic particles (SEPs). The electromagnetic output consists of the quasi-steady black body radiation with the superposition of flare emission. The mass and electromagnetic emissions are often coupled: flares and CMEs represent two different manifestations of the energy release from solar source regions (e.g., Asai et al. 2013). SEPs are accelerated in fast CME-driven shocks as well as by flare reconnection (see e.g., Reames 1999, 2013). The source regions of flares and CMEs are closed magnetic field regions such as active regions and filaments (see e.g., Srivastava et al. 2014). Active regions consist of sunspots of opposite polarity at the photospheric level. Filament regions do not have sunspots but consist of opposite polarity magnetic patches. The main difference between the two regions is the magnetic field strength: hundreds of gauss in the sunspot regions vs. tens or less gauss in filament regions. Sunspots also contribute to the variability in total solar irradiance (TSI). While sunspots decrease the TSI, the plages that surround the sunspots increase it, resulting in higher TSI when the Sun is more active.
The solar wind has fast and slow components. The fast component is of particular interest because it can compress the upstream slow solar wind forming a corotating interaction region (CIR). Coronal holes, the source of the fast solar wind, also exhibit remarkable variability in terms of their location on the Sun and size. Perhaps even more important than the CIR is the high-speed stream proper. It carries large nonlinear Alfvén waves, whose southward components cause reconnection at the magnetopause resulting in continuous sporadic plasmasheet injections into the nightside magnetosphere. These injections of anisotropic approximately 10- to 100-keV electrons cause the growth of an electromagnetic wave called ‘chorus’ and the chorus interacts with approximately 100-keV electrons accelerating them to MeV energies (Tsurutani et al. 2006, 2010; Thorne et al. 2013).
CMEs are launched into the solar wind, so the two mass outputs interact and exchange momentum affecting the propagation characteristics of CMEs in the interplanetary medium. CMEs also interact with the upstream heliospheric current sheet and other material left over from other injections. The variability manifested as solar flares, CMEs, SEPs, and high-speed solar wind streams directly affects space weather on short time scales. As noted above, all these phenomena are coupled not only near the Sun but throughout the inner heliosphere, including geospace and Earth’s ionosphere and atmosphere where the impact can be felt (Verkhoglyadova et al. 2014; Mannucci et al. 2014; Tsurutani et al. 2014).
This review covers the second phase of the CAWSES program, known as CAWSES-II, which began in 2009 and ended in 2013, roughly covering the rise to the maximum phase of solar cycle 24. There is strong evidence showing that solar cycle 24 is a relatively weak cycle (Tan 2011; Basu 2013). The birth of solar cycle 24 was remarkable in that the Sun emerged from an extremely deep minimum. The maximum phase of cycle 24 is of particular interest because the sunspot number was rather small (roughly half of the cycle 23 peak). The weak solar cycle resulted in a milder space weather, but there were other complications such as longer living space debris due to the reduced atmospheric drag. SCOSTEP conducted a year-long campaign known as ‘MiniMax24’ to document solar events and their geospace impact during the mild maximum phase of cycle 24. Additionally, the solar mid-term and long-term quasi-periodic cycles and their possible relationships with planetary motions from long-term observations of the relative sunspot number and microwave emission at frequency of 2.80 GHz were also investigated, and it was suggested that the mid-term solar cycles (periods <12 years) are closely related to the motions of the inner planets and of Jupiter (Tan and Cheng 2013).
This paper highlights some key results obtained on the variable phenomena in the Sun-Earth system during the CAWSES-II time frame. Detailed treatment of topics relevant to CAWSES-II can be found in the Living Reviews in Solar Physics: Hathaway (2010), Charbonneau (2010), Petrovay (2010), Chen (2011), Shibata and Magara (2011), Aschwanden (2011), Webb and Howard (2012), Usoskin (2013), Potgieter (2013), Lockwood (2013), Owens and Forsyth (2013).
The weak solar cycle 24 and its consequences
CAWSES-II focused on solar variability as the Sun approached the maximum of cycle 24. The rise phase of cycle 24 was already interesting because the Sun emerged from a deep solar minimum that gained particular interest among solar-terrestrial scientists (Selhorst et al. 2011; Tsurutani et al. 2011a; Dasso et al. 2012; Gopalswamy et al. 2012a; Solomon et al. 2013; Lean et al. 2014; Potgieter et al. 2014). Solar activity is typically represented by the international sunspot number (SSN), but there are many other measures, which are needed for a complete understanding of the solar variability. In particular, measurements of the polar field strength, tilt angle of the heliospheric current sheet (Smith et al. 1978), latitudes of filament locations, and coronal streamers all provide complementary information on the solar activity as the Sun climbed toward its weak maximum around 2013.
Solar polar field strength
Implications for the solar dynamo
According to the Babcock-Leighton mechanism of the solar dynamo, the polar field strength of one cycle determines the strength of the next cycle. The so-called polar precursor method of predicting the strength of a solar cycle using the peak polar field strength of the preceding minimum has been fairly accurate (see e.g., Svalgaard et al. 2005; Jiang et al. 2013a; Muñoz-Jaramillo et al. 2013; Zolotova and Ponyavin 2013). Recent discussion on the precursor method can be found in Petrovay (2010) and Pesnell (2014) among others. In addition to the traditional polar field measurements, proxies such as H-alpha synoptic charts (Obridko and Shelting 2008) and the polar microwave Tb (Gopalswamy et al. 2012a) can also be used to predict the strength of the activity cycle. The polar microwave Tb is exceptionally good because it is highly correlated (correlation coefficient r = 0.86) with the polar field strength: B = 0.0067 Tb - 70 G (Gopalswamy et al. 2012a).
How do we understand the weak cycle 24? Jiang et al. (2013b) considered several possibilities such as (i) the accuracy of SSN, (ii) sunspot tilt angle variation, and (iii) the variation in the meridional circulation during cycle 23. They were able to reproduce the lower polar field during the cycle 23/24 minimum using a 55% increase of the meridional flow in their model. They also found that a 28% decrease of the mean tilt angle of sunspot groups can explain the low polar field, but this would not be consistent with the observed time of polar field reversals. They concluded that the nonlinearities in the polar field source parameters and in the transport parameters play important roles in the modulation of the polar field.
Implications for the long-term behavior of the Sun
The Sun is known to have variability on time scales up to millennia (see Usoskin 2013 for a review). One obvious question is whether the weakening of the activity observed in cycle 24 will continue further. Javaraiah (2015) examined the north-south asymmetry of sunspot areas binned into 10° latitudes and examined various periodicities. They found periodicities of 12 and 9 years, respectively, during low-activity (1890 to 1939) and high-activity (1940 to 1980) intervals. They also inferred that cycle 25 may be weaker than cycle 24 by approximately 31%. Several authors have discussed the possibility of a global minimum over the next several cycles (see e.g., Russell et al. 2013a; Lockwood et al. 2011; Steinhilber and Beer 2013; Zolotova and Ponyavin 2014; Ruzmaikin and Feynman 2014). Zolotova and Ponyavin (2014) reported that the protracted cycle 23 is similar to the cycles immediately preceding the Dalton and Gleissberg-Gnevyshev minima, suggesting that the Sun is heading toward such a grand minimum.
But the most important is that the diminished solar activity has immediate consequences for the society. When the Maunder Minimum occurred in the late 1600s, the technology was not seriously affected by the Sun. Today’s technology is extensively coupled to solar activity, so the effect is readily recognized. For example, the weak solar activity has resulted in reduced atmospheric drag on satellites increasing their lifetime. On the other hand, space debris do not burn up quickly, thus posing additional danger to the operating satellites. The geomagnetic disturbances have been extremely mild, with the weakest level of geomagnetic storms since the space age.
The weakest geomagnetic activity on record: cycle 23 minimum
The figure shows that cycle 23 extended from 1996 to 2008 and is the longest in the space era (12.6 years). For comparison, the length of solar cycles 20 through 22 was 11.7, 10.3, and 9.7 years, respectively. The values below the red lines have been shaded for emphases (in the case of cosmic rays, the values above the red line are shaded). It can be noted that the Bo, Vsw, and ap index values for the cycle 23 minimum are considerably lower than the cycle 22 minimum values. The minimum in ap is broad and extends from day 97, 2008, until day 95, 2010. The onset and end times are somewhat arbitrary. There is a minimum geomagnetic activity interval in cycle 22 (day 106, 1996, to day 23, 1998).
It is surmised that nothing has changed on the speed of the high-speed streams emanating from coronal holes during solar minimum. The terminal speed is still approximately 750 to 800 km/s. However, this is the speed for the central portion of the hole. As the high speed stream expands into interplanetary space, it does not simply propagate radially outward but expands into nearby space, leading to ‘super-radial’ expansion as shown in the schematic of Figure 5. At the sides of the high-speed stream, the speed and the amplitude of the entrained Alfvén waves will be reduced. This is the portion of the high-speed streams that hit the Earth’s magnetosphere.
Thus, the low solar magnetic fields, the lack of CMEs, and the midlatitude location of small coronal holes all contribute to the all-time minimum in the geomagnetic activity between 2008 and 2009. It is noted that in Figure 3, a similar feature can be noted in the cycle 22 minimum, but the feature is less prominent.
Coronal mass ejections and flares
Origin of solar eruptions
Although it is well established that CMEs and their interplanetary manifestations, ICMEs, and flares originate from closed-field regions on the Sun such as active regions and filament regions, the current level of understanding is not sufficient to predict when an eruption might occur in such a region. Two basic processes seem to be involved: energy storage and triggering. The energy storage can be identified from non-potentiality of the source region such as magnetic shear or accumulated helicity (Tsurutani et al. 2009a; Kazachenko et al. 2012). Zhang et al. (2012a) studied the magnetic helicity of axisymmetric power-law force-free fields and focused on a family whose surface flux distributions are defined by self-similar force-free fields. The results suggest that there may be an absolute upper bound on the total magnetic helicity of all bipolar axisymmetric force-free fields.
In addition to the energy storage, a trigger in the form of a magnetic disturbance seems to be necessary, which causes a pre-eruption reconnection (Kusano et al. 2012). These authors suggest that observing these triggers is important for predicting eruptions and that we can predict eruptions only by a few hours before the eruption. For longer-term predictions, one has to resort to probabilistic methods. Huang et al. (2011) presented a study of a coronal mass ejection (CME) with high temporal cadence observations in radio and extreme-ultraviolet (EUV). The radio observations combined imaging of the low corona with radio spectra in the outer corona and interplanetary space. They found that the CME initiation phase was characterized by emissions that were signatures of the reconnection of the outer part of the erupting configuration with surrounding magnetic fields. Later on, a main source of emission was located in the core of the active region, which is an indirect signature of the magnetic reconnection occurring behind the erupting flux rope. Energetic particles were also injected in the flux rope and the corresponding radio sources were detected. Other radio sources, located in front of the EUV bright front, traced the interaction of the flux rope with the surrounding fields. They found that imaging radio emissions in the metric range can trace the extent and orientation of the flux rope which was later detected in interplanetary space.
Long-term behavior of CME rates
Importance of CMEs for space weather
For space weather effects, more energetic CMEs need to be examined. Gopalswamy et al. (2014b) started with flares of soft X-ray size ≥C3.0. This criterion avoids the effect of soft X-ray background level and its variability between the two cycles. For example, approximately 20% of flares of size <C1 cannot be detected by GOES during cycle 23, while the corresponding fraction is 6% for cycle 24. However, for flares of size ≥C3.0, no flares go undetected. The source locations of flares were obtained from the online Solar Geophysical Data (SGD) or identified using near-surface observations from a number of sources: EUV images from SOHO, STEREO (Solar Terrestrial Relations Observatory), and SDO (Solar Dynamics Observatory), soft X-ray images from Yohkoh, microwave images from the Nobeyama Radioheliograph, and H-alpha images from various observatories. For each of these flares, the association of a CME was checked using the SOHO/LASCO catalog (http://cdaw.gsfc.nasa.gov, Gopalswamy et al. 2009a) to compile the properties of the associated CMEs. Only flares that occurred within 30° from the limb were considered, so that the speed and width measurements of CMEs were subject to minimal projection effects.
CME speed and width distributions
Comparison between CME numbers in solar cycles 23 and 24
Cycle 23 a
W < 30°
W ≥ 30°
W ≥ 60°
W = 360°
V ≥ 900 km/s and W ≥ 60°
≥C3.0 flares, limb
Anomalous expansion of CMEs in cycle 24
CME mass distribution in cycles 23 and 24
The weak state of the heliosphere
Despite large international efforts in understanding FDs, there is still a lot to learn. The current model of FDs consisting of two-step decrease has recently been questioned. It is not clear if only a subset of CMEs originating from the disk center is effective in causing FDs (Jordanova et al. 2012). However, the study of FDs has been gaining interest in recent times because of the space weather applications. For example, the development of Global Muon Detector Network (GMDN - Munakata et al. 2005; Fushishita et al. 2010; Rockenbach et al. 2011) has greatly enhanced the possibility of forecasting ICME arrival using the network (see e.g., Rockenbach et al. 2014 for a review).
Spatial structure of CMEs
Even before the discovery of white-light CMEs, the concept of magnetic loops from the Sun driving shocks was considered (Gold 1962). In Gold’s picture, a magnetic bottle from the Sun drives a fast magnetosonic shock which stands at certain distance from the bottle. Such a shock was first identified by the Mariner 2 mission in 1962 (Sonett et al. 1964). Koomen et al. (1974) identified white-light CMEs with the Gold bottle. Burlaga et al. (1981) confirmed the basic picture of Gold using in situ data by identifying the shock, sheath, and the driving magnetic structure. Near the Sun, MHD shocks were inferred from metric type II radio bursts for several decades ago (see e.g., Nelson and Melrose 1985). The overall CME structure consisting of a flux rope enclosing a prominence core and driving a shock outside has been considered by theorists a while ago (e.g., Kuin and Martens 1986), but it took another two decades before the white-light shock structure of CMEs was observed in coronagraphic images (Sheeley et al. 2000). A recent study based on coronagraph observations concluded that a flux rope structure can be discerned in approximately 40% of CMEs observed near the Sun (Vourlidas et al. 2013).
White-light and EUV signatures of CME-driven shocks
Fast forward interplanetary shocks (hereafter simply called ‘shocks’) are driven by either fast CMEs or high-speed streams. So far, no ‘blast wave’ shocks have been detected in the interplanetary medium by spacecraft instrumentation. Shocks compress and heat the upstream plasma and magnetic fields (Kennel et al. 1985). Thus, the immediate downstream (or sheath) region may be visible at times. Shocks form from a steepening of magnetosonic waves. To identify whether a wave is a shock or not, it must be shown to have a supermagnetosonic speed in its normal direction. Methods of analyses can be found in Tsurutani and Lin (1985) and the geoeffectiveness of shocks and discontinuities in Tsurutani et al. (2011b).
There have been several recent studies on white-light shocks (Vourlidas et al. 2003; Michalek et al. 2007; Gopalswamy et al. 2008a; Gopalswamy et al. 2009; Gopalswamy et al. 2009b; Ontiveros and Vourlidas 2009; Bemporad and Mancuso 2011; Gopalswamy and Yashiro 2011; Maloney and Gallagher 2011; Kim et al. 2012; Poomvises et al. 2012) that have provided a better understanding of the CME structure beyond the classical three-part structure (Hundhausen 1987). The availability of STEREO and SDO observations increased our ability to visualize the CME-shock system and understand the shock formation and coronal plasma properties.
The dome structure surrounding newly erupted CMEs has been recognized as the three-dimensional counterpart of the so-called EIT waves (Patsourakos and Vourlidas 2009; Veronig et al. 2010; Ma et al. 2011; Kozarev et al. 2011; Warmuth and Mann 2011; Gallagher and Long 2011; Harra et al. 2011; Gopalswamy et al. 2012b; Selwa et al. 2013; Temmer et al. 2013; Liu and Ofman 2014; Nitta et al. 2013a). The wave nature of EUV waves was also established based on the fact that they are reflected from nearby coronal holes (Long et al. 2008, 2013; Gopalswamy et al. 2009c; Olmedo et al. 2012; Shen et al. 2013a; Kienreich et al. 2013: Kwon et al. 2013). Gopalswamy and Yashiro (2011) estimated the coronal magnetic field within the SOHO coronagraphic field of view (6 to 23 Rs) using the fact that the standoff distance of the white-light shock with respect to the radius of curvature of the driving flux rope is related to the shock Mach number and the adiabatic index (Russell and Mulligan 2002; Savani et al. 2012). Since the shock speed is measured from the coronagraphic images, these authors were able to derive the Alfvén speed and magnetic field in the ambient medium. Poomvises et al. (2012) extended this technique to the interplanetary medium and showed that the derived magnetic field strength is consistent with the HELIOS in situ observations. This technique will be extremely important to compare future in situ observations from missions to the Sun such as Solar Orbiter and Solar Probe Plus, currently under development (Müller et al. 2013). The standoff-distance technique was also applied to a CME-shock structure observed by SDO on 13 June 2010, which showed that the technique can work as close to the Sun as 1.20 Rs, where the shock first formed (Gopalswamy et al. 2012b; Downs et al. 2012). The shock formation heights derived from SDO/AIA and STEREO/EUVI have provided direct confirmation that CME-driven shocks have enough time to accelerate particles to GeV energies from a height of approximately 1.5 Rs before they are released when the CME reaches a height of about 3 to 4 Rs (Gopalswamy et al. 2013a,b; Thakur et al. 2014). The low shock formation height applies only to those CMEs, which quickly accelerate and attain high speeds (see e.g., Bein et al. 2011).
Shocks inferred from radio observations
Type II radio bursts in the metric domain, traditionally observed from ground-based observatories, indicate shock formation very close to the Sun (e.g., Kozarev et al. 2011; Ma et al. 2011; Gopalswamy et al. 2012b). Imaging these bursts provides important information such as the magnetic field in the ambient medium (Hariharan et al. 2014). These bursts indicate the height of shock formation in the corona as evidenced by EUV shocks and Moreton waves (see e.g., Asai et al. 2012a, 2012b). Radio emission from interplanetary shocks in the form of type II bursts provides important information of shock propagation in the heliosphere (Gopalswamy 2011). CMEs with continued acceleration beyond the coronagraph field of view (FOV) may form shocks at large distances where they become super-magnetosonic (faster than the upstream magnetosonic wave speed). Shocks forming at large distances of the Sun may or may not produce type II radio bursts (Gopalswamy et al. 2010c). Radio-quiet CMEs (those lacking type II radio bursts) typically have positive acceleration in the coronagraphic field of view and become super-magnetosonic in the interplanetary (IP) medium at large heliocentric distances. Deceleration of radio-loud CMEs near the Sun and the continued acceleration of radio-quiet CMEs into the IP medium make them appear similar at 1 AU. However, there is a better chance that radio-loud CMEs produce an energetic storm particle event (Mäkelä et al. 2011) and strong sudden commencement/sudden impulse (Veenadhari et al. 2012), suggesting that stronger shocks near the Sun do matter. In fact Vainio et al. (2014) have shown that the cut-off momentum of particles observed at 1 AU can be used to infer properties of the foreshock and the resulting energetic storm particle event, when the shock is still near the Sun.
By combining STEREO/HI (heliospheric imager) observations, interplanetary radio bursts observations, and in situ measurements from multiple vantage points, Liu et al. (2013) showed that it is possible to track CMEs and shocks. In particular, they were able to study CME interaction signatures in the radio dynamic spectrum. The drift rate of the type II radio bursts can also be converted into shock speed for comparison with the CME speed derived from HI observations, providing a method to predict shock arrival (e.g., Xie et al. 2013a).
Shock critical Mach numbers
There is renewed interest in shock critical Mach numbers and their evolution with heliocentric distance (Gopalswamy et al. 2012b; Bemporad and Mancuso 2011, 2013; Vink and Yamazaki 2014). For example, some radio-loud shocks may dissipate before reaching 1 AU (Gopalswamy et al. 2012c) indicating that the Mach number is dropping to 1 or below. Bemporad and Mancuso (2011) concluded that the supercritical region occupies a larger surface of the shock early on but shrinks to the nose part of the shock as it travels away from the Sun. Vink and Yamazaki (2014) introduced a different critical Mach number (Macc), which is substantially larger than the first critical Mach number (Mcrit) of quasi-parallel shocks (Kennel et al. 1985) but similar to Mcrit of quasi-perpendicular shocks. According to these authors, the condition Macc > √5 seems to be required for particle acceleration, which may be relaxed when seed particles exist.
The clear identification of an interplanetary magnetic cloud (MC) with a CME by Burlaga et al. (1982) replaced the Gold bottle by a flux rope. Although the MC definition by Burlaga et al. (1982) was narrower than the flux rope definition (magnetic field twisted around an axis), the terms ‘flux rope’ and ‘MC’ are interchangeably used after Goldstein (1983) showed that the MC magnetic field can be modeled by a flux rope. All theories of CME eruption and propagation use the flux rope as the fundamental structure in their calculations, either preexisting or formed during eruption (Yeh 1995; Chen et al. 1997; Riley et al. 2006; Forbes et al. 2006; Chen 2012; Kleimann 2012; Lionello et al. 2013; Janvier et al. 2013; Démoulin 2014 Extensive CME observations from the SOHO mission have helped perform many studies on CME flux ropes. Chen et al. (1997) showed that the observed CME structure in the LASCO field of view can be interpreted as the two-dimensional projection of a three-dimensional magnetic flux rope with its legs connected to the Sun.
The CME flux rope is thought to be either pre-existing or formed out of reconnection during the eruption process and is observed as an MC in the interplanetary medium (see e.g., Gosling 1990; Leamon et al. 2004; Qiu et al. 2007). On the other hand, it is possible that a set of loops from an active region on the Sun can simply expand into the interplanetary (IP) medium and can be detected as an enhancement in the magnetic field with respect to the ambient medium (Gosling 1990) without any flux-rope structure. The in situ magnetic signatures will be different in the two cases. A spacecraft passing through the flux rope center will see a large, smooth rotation of the magnetic field throughout the body of the interplanetary CME (ICME), while the expanded loop system will show no rotation. If we take just the IP observations, we may be able to explain MCs as flux ropes and non-MCs as expanding loops. However, they should show different charge-state characteristics (see e.g., Aguilar-Rodriguez et al. 2006; Gopalswamy et al. 2013c) because of the different solar origins. The flux rope forms during the flare process and hence is accessed by the hot plasma resulting in high charge states inside MCs when observed at 1 AU. Expanding loops, on the other hand, should not have high charge states because no reconnection is involved. Riley and Richardson (2013) analyzed Ulysses spacecraft measurements of ICMEs and concluded the ICME may not appear as MCs because of observing limitations or the initiation mechanism at the Sun may not produce MCs.
In a series of two coordinated data analysis workshops (CDAWs), a set of structure of CMEs, 54 CME-ICME pairs were analyzed to study the flux-rope nature of CMEs (see Gopalswamy et al. 2013d for the list of papers based on these CDAWs). It was found that MCs and non-MCs were indistinguishable based on their near-Sun manifestations such as white-light CMEs and flare post-eruption arcades. In particular, the CMEs were fast and the flare arcades were well defined (Yashiro et al. 2013). Fe and O charge states at 1 AU were also indistinguishable between MCs and non-MCs, suggesting a similar eruption mechanism for both types at the Sun (Gopalswamy et al. 2013c). Combined with the fact that CMEs can be deflected toward or away from the Sun-Earth line (Gopalswamy et al., 2009d), the observing geometry (i.e., the observing spacecraft may not cross the flux rope axis) seems to be the primary reason for the non-MC appearance of flux ropes (see e.g., Kim et al. 2013). Many authors have advocated that all ICMEs are flux ropes (Marubashi 2000; Owens et al. 2005; Gopalswamy 2006a), but the single point observations at 1 AU may miss it. Marubashi et al. (2015) showed that almost all ICMEs can be fit to a flux rope if a locally toroidal flux rope model is considered in addition to the cylindrical flux rope model. Similarly, the active region helicity and the helicity of the ICMEs were in good agreement (Cho et al. 2013).
Using SDO/AIA data, Zhang et al. (2012b) reported that flux ropes exist as a hot channel before and during an eruption. The structure initially appeared as a twisted and writhed sigmoid with a temperature as high as 10 MK and then transformed into a semi-circular structure during the slow-rise phase, which was followed by a fast acceleration and flare onset. Cheng et al. (2013) reported that the hot channel rises before the first appearance of the CME leading front and the flare onset of the associated flare. These results indicate that the hot channel acts as a continuous driver of the CME formation and eruption in the early acceleration phase. Li and Zhang (2013a, 2013b) reported on the eruption of two flux ropes from the same active region within 25 minutes of each other on 23 January 2012. The two flux ropes initially rose rapidly, slowed down, and accelerated again to become CMEs in the coronagraph FOV. The two CMEs were found to be interacting in the coronagraph FOV as observed by SOHO and STEREO (Joshi et al. 2013). Li and Zhang (2013c) also studied homologous flux ropes from active region 11745 during 20 to 22 May 2013. All flux ropes involved in the eruption had a similar morphology.
Shock normal angles and particle acceleration
It has been shown by detailed studies that the normal of the shocks relative to the upstream magnetic field is important for the efficiency of particle acceleration. The physical reasoning is that for quasi-parallel shocks where the normal and upstream field are nearly aligned, shock-reflected particles create upstream waves by beaming instabilities (Tsurutani and Rodriguez, 1981; Tsurutani et al. 1983b). The upstream (and downstream) waves act as two walls of a Fermi accelerator (Lee, 1983) leading to exceptional particle acceleration. Evidence of this has been demonstrated by Kennel et al. 1984a, 1984b). This microscale structure was not examined in the particle acceleration studies discussed in earlier sections.
Propagation effects: deflection, interaction, and rotation of CMEs
Once a CME is ejected from the Sun, its 3D geometry at a far-away location such as Earth depends on the evolution in the changing background solar wind and magnetic field (Temmer et al. 2011). The CME flux ropes expand, so the magnetic content typically decreases. CMEs can be deflected in the latitudinal and longitudinal directions by pressure gradients (magnetic + plasma). The CME flux ropes can also be distorted by changing flow speeds in the background. Finally, CMEs may also rotate, so the orientation of the magnetic field inferred from solar observations may not match what is observed at Earth. It is possible that many of these effects can occur simultaneously (Nieves-Chinchilla et al. 2012). At present, there are a few techniques to connect CMEs observed at the Sun with their interplanetary counterparts. Interplanetary type II bursts detected by Wind/WAVES and STEREO/WAVES instruments can track CME-driven shocks all the way from the Sun to the observing spacecraft located at 1 AU (Xie et al. 2013a). Interplanetary scintillation (IPS) observations track turbulence regions surrounding CMEs (typically the sheath region) also over the Sun-Earth distance (see e.g., Manoharan 2010; Jackson et al. 2013). The heliospheric imagers on board STEREO track CMEs in white light over the Sun-Earth distance (e.g. Möstl and Davies 2013).
A combination of CME tracking in the inner heliosphere using STEREO heliospheric imagers (HIs) and numerical simulations have greatly enhanced CME propagation studies. The interaction of CMEs with the ambient medium is the primary propagation effect. This interaction is represented by an aerodynamic drag that dominates beyond the coronagraphic field of view. Close to the Sun, the propelling force and gravity dominate (see e.g., Vrsnak and Gopalswamy 2002). Defining the background is one of the key inputs needed for understanding CME propagation (see e.g., Roussev et al. 2012; Arge et al. 2013). However, there are other processes that can significantly affect the propagation of CMEs: CME-CME interaction (see e.g., Gopalswamy et al. 2001a; 2012c; Temmer et al. 2012; Harrison et al. 2012; Lugaz et al. 2012; Liu et al. 2014a; Sterling et al. 2014; Temmer et al. 2014) and CME deflection by large-scale structures such as coronal holes and streamers (Gopalswamy 2010, 2009d; Shen et al. 2011; Gui et al. 2011; Wood et al. 2012; Kay et al. 2013; Panasenco et al. 2013; Gopalswamy and Mäkelä 2014).
Different types of interaction become predominant during different phases of the solar cycle. During the rise phase, when polar coronal holes are strong and CMEs originate at relatively higher latitudes, the polar coronal holes are effective in deflecting CMEs (e.g., Gopalswamy et al. 2008b). During the maximum phase, CMEs occur in great numbers, so CME-CME interaction is highly likely (e.g., Gopalswamy et al. 2012c; Lugaz et al. 2013; Chatterjee and Fan 2013; Farrugia et al. 2013; Kahler and Vourlidas 2014). CME interactions can also result in CME deflection and merger (Shen et al. 2012). In the declining phase, low-latitude coronal holes appear frequently, so CME deflection by such coronal holes becomes important (Gopalswamy et al. 2009d; Mohamed et al. 2012; Mäkelä et al. 2013). The deflections are thought to be caused by the magnetic pressure gradient between the eruption region and the coronal hole (Gopalswamy et al. 2010d; Shen et al. 2011; Gui et al. 2011).
Determining the initial orientation of flux ropes has been possible by fitting a flux rope to coronagraph observations (Thernisien 2011; Xie et al. 2013b). Such fitting already provides a lot of information on the deviation of CME propagation direction from the radial (e.g., Gopalswamy et al. 2014b). Isavnin et al. (2014) defined such initial flux rope orientation using extreme-ultraviolet observations of post-eruption arcades and/or eruptive prominences and coronagraph observations. Then, they propagated the flux rope to 1 AU in a MHD-simulated background solar wind and used in situ observations to check the results at 1 AU. They confirmed that the flux-rope deflection occurs predominantly within 30 Rs, but a significant amount of deflection and rotation happens between 30 Rs and 1 AU. They also found that that slow flux ropes tend to align with the streams of slow solar wind in the inner heliosphere.
CME arrival at Earth
There have been many attempts recently to convert the knowledge gained on CME propagation to predict the arrival times at 1 AU. The CME travel time essentially depends on the accurate estimate of the space speed of CMEs and the background solar wind speed; these have been estimated based on single view (SOHO) as well as from multiple views (SOHO and STEREO). Davis et al. (2010) found that CME speeds derived from STEREO/COR2 and Thernisien (2011) forward-fitting model were in good agreement, although CME speeds changed in the HI FOV depending on the near-Sun speed. Millward et al. (2013) developed the CME analysis tool (CAT), which models CMEs to have a lemniscate shape, which is similar to ice cream cone model. They showed that the leading-edge height and half-angular width of CMEs can be determined more accurately using multi-view data. Colaninno et al. (2013) tracked nine CMEs continuously from the Sun to near Earth in SOHO and STEREO images and found that the time of arrival was within ±13 h. Gopalswamy et al. (2013e) considered a set of 20 Earth-directed halos viewed by SOHO and STEREO in quadrature, so as to obtain the true earthward speed of CMEs. When the speeds were input to the empirical shock arrival (ESA) model, they found that the ESA model predicts the CME travel time within about 7.3 h, which is similar to the predictions by the ENLIL model. They also found that CME-CME and CME-coronal hole interaction can lead to large deviations from model predictions. Vrsnak et al. (2014) compared the arrival-time predictions from the ‘WSA-ENLIL + cone model’ and the analytical ‘drag-based model’ (Vršnak et al. 2013) and found that the difference in predictions had an absolute average of 7.1 h. Compared with observations, the drag-based model had an average absolute difference of 14.8 h, similar to that for the ENLIL model (14.1 h). Xie et al. (2013a) compared travel times of CMEs when ENLIL + cone model and ENLIL + flux rope model were used. They found that the ENLIL + flux rope model results showed a slight improvement (4.9 h vs. 5.5 h). They also found that predictions based on kilometric type II bursts improved significantly when the ENLIL model density was used rather than the average solar wind plasma density in deriving shock speeds from the type II drift rate. The improvement was typically better by approximately 2 h. Möstl et al. (2014) derived the absolute difference between predicted and observed ICME arrival times for 22 CMEs as 8.1 h (rms value of 10.9 h). Empirical corrections to the predictions reduced arrival times to within 6.1 h (rms value of 7.9 h). Echer et al. (2010) attempted to identify the solar origins of the November 2004 superstorms on Earth using existing interplanetary propagation routines published in the literature. They found that during highly active solar intervals, the predictions were sometimes ambiguous, in agreement with the comments above. Thus, there has been a steady progress in predicting the arrival time of CMEs, which need to be continued and expended to the prediction of Bz values, which are crucial to predict the strength of geomagnetic storms.
CMEs and geomagnetic storms
One of the direct consequences of CMEs arriving at Earth’s magnetosphere is the geomagnetic storm. The primary link between a geomagnetic storm and a CME is the out of the ecliptic component (Bz) of the interplanetary magnetic field (Gonzalez et al., 1994; Zhang et al. 2007; Gopalswamy 2008; Echer et al. 2008a, 2008b, 2013; Cid et al. 2012). Echer et al. (2008a) conclusively showed that for all 90 major (Dst < −100 nT) storms that occurred during cycle 23, it was the Bz component that was responsible for the storms (some people have thought that it was possible that the IMF By component was also important). When Bz is negative (south pointing), then the CME field reconnects with Earth’s magnetic field (Dungey 1961) causing the geomagnetic storm. While the Bz component is negligible in the quiet solar wind, CMEs contain Bz by virtue of their flux rope structure. Fast CMEs drive shocks, so the compressed sheath field between the flux rope and the shock can also contain Bz (Tsurutani et al. 1988). Thus, both the flux rope and sheath can be the source of Bz and hence cause geomagnetic storms. One of the common indicators of the strength of geomagnetic storms is the Dst index (expressed in nT), which is computed as the horizontal component of Earth’s magnetic field measured at several equatorial stations (now, a SYM-H index is available which is essentially a 1-min resolution Dst index). Yakovchouk et al. (2012) reported significant difference between the local and global peak storm intensities: the local storm minima were found to be 25% to 30% stronger than the global minima. Here, we consider only global peak intensities. Major storms have Dst ≤ −100 nT and are mostly caused by CMEs. Echer et al. (2008a) studied all 90 storms with Dst ≤ −100 nT for cycle 23. They divided ICMEs proper from their upstream sheaths. They found that roughly half of the storms that were caused by CME/sheaths were due to CMEs and half due to sheath fields. Yakovchouk et al. (2012) found that 10% of major storms are caused by CIRs. Echer et al. (2008a) determined that 13% of the cycle 23 storms were caused by CIRs.
It should be noted that all studies of ‘superstorms’ or storms with Dst intensities < −250 nT have been caused by magnetic clouds (Tsurutani et al. 1992a; Echer et al. 2008b). Such intense storms are not caused by sheaths or CIRs. One might ask why not? The argument is a simple one. The slow solar wind magnetic field is approximately 5 to 7 nT. It has been shown by Kennel et al. (1985) that fast shocks can compress the magnetic field by a maximum factor of approximately 4, regardless of the shock Mach number. Thus, interplanetary sheath fields and CIR magnetic fields should have maximum field strengths of approximately 20 to 35 nT. In contrast, magnetic cloud fields have been 50 to 60 nT and in exceptional cases approximately 100 nT. Thus, even if the sheath and CIR magnetic fields are totally southward, they are small by comparison to MC fields.
There is one exception to this above explanation. In cases of ARs where there are multiple CMEs and multiple shocks, the shocks can ‘pump up’ sheath magnetic field intensities. This has been shown to be the case for both CAWSES intervals of study (Tsurutani et al. 2008; Echer et al. 2011c.
Major geomagnetic storms of cycle 24 (Dst < −100 nT)
Date and time of storm
V a km/s
W b degree
Bz (nT) location c
Six of the 11 CME storms in Table 2 were due to southward Bz in the ICME sheaths. These include one event in which the sheath was actually a preceding magnetic cloud: the shock from a disk-center CME on 16 Feb 2014 from S11E01 associated with a M1.1 flare at 09:20 UT entered into a preceding fully south (FS) cloud. The source of the magnetic cloud itself is not clear, but most likely a faint CME associated with an eruption in the northwest quadrant on 14 Feb 2014 around 02:33 UT (a faint CME at 4:28 UT CME was barely discernible in LASCO images). The shock compressed the preceding CME and enhanced the Bz that caused the storm. This is a good example that ICMEs can be affected by shocks from other eruptions. In the remaining five storms, southward Bz was in the cloud portion. Even cycles such as cycle 24 are supposed to have more north-south (NS) clouds (the ones with leading northward Bz). However, there was only one such cloud in Table 2. All others were of FS clouds, which are high-inclination clouds with south-pointing axial field. Note that the Bz values ranged from −11.1 to −28.7 nT with an average of −18.5 nT. The Bz values in cycle 24 were in a narrower range compared to those in cycle 23.
Flares and the ionosphere
It has long been known that solar flares create sudden ionospheric disturbances or SIDs (Thome and Wagner 1971; Mitra 1974). However, the extreme intensity of the Halloween flares and the rise of using global positioning systems (GPS) for ionospheric research has allowed major advances to be made in flare ionospheric research (Tsurutani et al. 2005; Afraimovich et al. 2009). Now with ground-based receivers virtually everywhere on Earth (with ocean coverage still a bit of a problem), high-time resolution, global coverage is now possible.
The solar flare causes the largest TEC enhancement at the subsolar region with a TEC enhancement of 22 TEC units. The nightside region shows no TEC change, as expected. This is the largest ionospheric TEC change due to a solar flare ever detected.
CMEs and the ionosphere
Figure 19 shows the ‘dayside superfountain effect’ for the 30 October 2003 Halloween magnetic storm. The CHAMP satellite passes before the storm (blue trace), which shows the two equatorial ionization anomalies (EIAs) located at approximately ±10°. With time, the ionosphere and EIAs are uplifted to higher magnetic latitudes and have higher intensities. In the first pass after storm onset (red curve), the EIAs have peak intensities of approximately 200 TECU at approximately ±20° MLAT. In the following pass, a peak intensity of approximately 330 TECU is detected at approximately 30° MLAT. The cause of this remarkable feature is the interplanetary dawn-dusk electric field which uplifts the upper ionosphere by E × B convection (the Earth’s magnetic field is aligned in a north-south direction at the magnetic equator). As the electrons and ions are convected to higher altitudes and latitudes, solar irradiation replaces the uplifted plasma by photoionization, leading to an overall increase in the TEC (Tsurutani et al. 2004; Mannucci et al. 2005).
Coronal holes and CIRs
Coronal holes play a number of important roles in Sun-Earth connection. Polar coronal holes indicate the strength of the polar field and hence the level of solar activity in the following cycle (Gopalswamy et al. 2012a; Selhorst et al. 2011; Shibasaki 2013; Mordvinov and Yazev 2014; Altrock 2014). Coronal holes in the equatorial region are good indicators of imminent high-speed streams (HSS) and CIRs arriving at Earth (Tsurutani et al. 1995, 2006; Cranmer 2009; Verbanac et al. 2011; Akiyama et al. 2013; Borovsky and Denton 2013). The empirical relationships established between HSS characteristics and the related geomagnetic activity provides an advance warning of impending CIR storms (Tsurutani et al. 2006; Verbanac et al. 2011). Coronal holes also deflect CME-driven shocks and CMEs that have important space weather consequences (Gopalswamy 2010, 2009d; Olmedo et al. 2012; Kay et al. 2013; Mäkelä et al. 2013). The deflection by coronal holes can be so large that CMEs originating from close the disk center of the Sun do not arrive at Earth while the shocks do. Coronal-hole deflection may also make the shock and the driving flux rope appear unaligned (Wood et al. 2012). Both high-speed streams and CIRs result in various types of magnetospheric responses (Tsurutani et al. 2006; Denton and Borovsky 2012; Borovsky and Denton 2013). Coronal holes also seem to play a critical role in deciding whether CMEs originating at latitudes >30° can produce ground-level enhancement events (Gopalswamy and Mäkelä 2014), although their effect on large SEP events has not been conclusive (Kahler et al. 2014). In this section, we provide an overview of the recent progress on the geomagnetic response.
CIRs, CIR storms, and HSS geomagnetic activity
High-speed streams (HSSs) originate from coronal holes and form CIRs when they collide with the slower solar wind ahead (see e.g., Smith and Wolf 1976; Gosling 1996). The compressed interaction region has a higher density and temperature, and the magnetic field intensities and fluctuations are amplified. When the field within the CIR contains southward Bz, geomagnetic activity ensues (Borovsky 2013). However, the fields within CIRs are typically highly fluctuating (compressed Alfvén waves). Thus, the character is considerably different from that of the magnetic fields within magnetic clouds. CIRs typically do not cause magnetic storms with Dst < −100 nT (Tsurutani et al. 1995, 2006; Echer et al. 2008a). The geomagnetic activity is typically in the range −50 nT < Dst < −100 nT. Geomagnetic storms caused by CMEs and CIRs (plus the following HSSs) differ in some important ways (Denton et al. 2006; Jordanova et al. 2012; Liemohn and Katus 2012; Verbanac et al. 2013; Borovsky and Denton 2013; Keesee et al. 2014). For example, the CMEs and CIRs (plus HSSs) have different magnetospheric responses resulting in different development of various current systems and geomagnetic activity within the Earth’s magnetosphere and ionosphere (Mannucci et al. 2005, 2008, 2012; Thayer et al. 2008; Lei et al. 2008, 2011; Verkhoglyadova et al. 2011, 2013, 2014; Verbanac et al. 2013). Keesee et al. (2014) performed superposed epoch analysis of 21 CME-driven and 15 CIR-driven storms during the June 2008 to April 2012 time frame and different evolutions of the ion temperature: the ion temperature increased in the recovery phase of CIR storms, while it increased rapidly at the onset of CME storms and cooled off during the main phase. Borovsky and Denton (2013) compared CIR storms associated with helmet streamers and pseudo-streamers. They found that pseudo-streamer CIR storms tend not to have a calm (Tsurutani et al. 1995) before the storm, with weaker superdense plasma sheet and electron radiation belt dropout.
Effect of the weak solar activity
Even though the CIRs were weak during the prolonged minimum (see e.g., Echer et al. 2011a), they had interesting effects on the ionosphere and atmosphere. The ionospheric response to the weak CIRs was marginal but observable (Araujo-Pradere et al. 2011). However, the weak recurrent geomagnetic activity due to CIRs did produce distinct variability in the thermospheric density at an altitude of 400 km above ground (Lei et al. 2011 and references therein). The thermosphere was found to respond globally with the density varying by approximately 75%. Most importantly, they were able to isolate the effect of geomagnetic activity from the EUV forcing because the EUV flux remained roughly constant during these CIR intervals. Thermospheric density variations also showed the periodicities in CIRs due to the spatial distribution of low-latitude coronal holes on the Sun, as did the ionosphere.
Figure 21 gives the solar cycle dependence of HILDCAAs. HILDCAAs are most often detected during the declining phase of the solar cycle, but they can be detected during the other phases as well.
The overall scenario of the magnetospheric relativistic electron acceleration starts at the Sun (Tsurutani et al. 2006, 2010). Supergranular circulation is the source of the Alfvén waves (Hollweg 2006). These Alfvén waves are carried from the coronal holes at the Sun to the Earth by the HSSs. The southward components of the Alfvén waves lead to magnetic reconnection and the geomagnetic activity indicated in the HILDCAAs. The injection of approximately 10- to 100-keV electrons by substorms/convection events within the HILDCAAs lead to chorus wave growth and the chorus accelerate electrons to relativistic energies. All features of this scenario have now been confirmed, except for the Alfvén wave source. Perhaps the Solar Probe Plus or Solar Orbiter will be able to finally confirm this last part of the scenario.
CIR and HSS ionospheric and atmospheric effects
Verkhoglyadova et al. (2011) studied ionospheric and atmospheric effects of HSSs in the solar minimum interval, 25 March to 26 April 2008. This was the study interval for the whole heliospheric interval (WHI) science team. As stated previously, this was near the solar minimum but about a year before the geomagnetic minimum.
Mlynczak et al. (2003) views NO irradiation as a ‘natural thermostat.’ Energy input into the upper atmosphere during high AE intervals is converted to heat and changes the distribution of NO in the thermosphere and its radiative properties. The radiation from NO cools the atmosphere. NO is believed to account for approximately 50% of the estimated energy input to the atmosphere from the magnetic storm.
The ionospheric TEC effects during CIRs/HSSs are weak but are clearly present (Verkhoglyadova et al. 2011). There is enhanced TEC both at low latitudes (second panel: −30° < MLAT < +30°) and at middle latitudes (third panel: 40° to 60° MLAT). The middle latitude disturbance is most likely due to particle precipitation in the auroral zone. However, the equatorial disturbances are less well understood. It may be associated with a disturbance dynamo effect or the dayside equatorial superfountain. More study is warranted.
Microwave enhancement in coronal holes and solar wind speed
Akiyama et al. (2013) also considered weak and strong CIRs but within solar cycle 23 (1996 to 2005). The weak and strong CIRs resulted in the Dst index > −100 nT and < −100 nT, respectively. The CIRs were associated with low-latitude coronal holes identified in EUV images from SOHO/EIT. They measured the area of the coronal hole in EUV and in microwaves (from the Nobeyama Radioheliograph). Coronal holes appear bright in microwaves, and the area of microwave enhancement overlaps with that in EUV but not completely. The area of microwave enhancement is typically a third of the area observed in EUV CH on the average (Gopalswamy et al. 2000). There are at least two reasons for this difference: (1) the microwave enhancement originates in the chromosphere, so the area is expected to be smaller due to magnetic field expansion, and (2) microwave enhancement is an indicator of higher magnetic field within the coronal hole, which is usually patchy within the coronal hole. A clear correlation between the CH size and the solar wind speed is well known (e.g., Nolte et al. 1976).
The lone major CIR storm of cycle 24
Large SEP events and GLE events
Solar energetic particles (SEPs) are part of the mass emission from the Sun, intimately connected to CMEs and flares. The current paradigm is that the material in the corona and IP medium is accelerated by CME-driven shocks to produce the large SEP events (e.g., Reames 1999, 2013). Detailed investigation on the connection between CME-driven shocks and SEPs became possible only after the advent of SOHO coronagraphs that routinely imaged CMEs in the coronal region from where the energetic particles are released. In particular, the highest energy (GeV) particles have been studied in relation to CMEs only for the past couple of solar cycles (Gopalswamy et al. 2012d; Nitta et al. 2012; Mewaldt et al. 2012; Miroshnichenko et al. 2013). Studies on the connection of SEPs to CMEs have been further enhanced by the extended coverage provided by STEREO in a number of ways, including observations all around the Sun. The heliospheric imagers and the inner coronagraph COR1 have extended the spatial domain over which CMEs are observed. In addition, EUV observations from STEREO/EUVI and SDO/AIA have improved our ability to study the early phase of CMEs (see e.g., Aschwanden et al. 2014). For example, the coronal height where shocks form can be readily determined from EUV and COR1 images, so that the time available for accelerating SEPs can be estimated accurately. Characterizing the size of large eruptions from behind the limb has been made possible by estimating the soft X-ray fluxes based on the known correlation between EUV flux and soft X-ray flux for frontside events (Nitta et al. 2013b).
Even though STEREO was launched by the end of 2006, no large SEP events occurred until the second half of 2010 in the CAWSES-II period. From then on, there were many SEP events that have been studied extensively. The STEREO particle detectors observed SEP events from off the Sun-Earth line and, together with detectors at Sun-Earth L1, provided information on the longitudinal distribution of SEP intensity (e.g., Dresing et al. 2012; Rouillard et al. 2012; Mewaldt et al. 2013). It also became possible to study the radial dependence of SEP intensities, thanks to the observations provided by MESSENGER spacecraft (Lario et al. 2013). Observations from multiple spacecrafts have also been used to test diffusive shock acceleration: Wang and Yan (2012) performed a dynamical Monte Carlo simulation of the CME-driven shock that occurred on 14 December 2006 using an anisotropic scattering law and found that the simulated results of the shock’s fine structure, particle injection, and energy spectrum were in good agreement with the observations.
Longitudinal dependence of SEP intensity
Lario et al. (2006, 2013) were able to fit a Gaussian to the ensemble of multispacecraft SEP peak intensities as a function of the longitudinal distance between the solar source (inferred from flare observations) and the magnetic connection point for each spacecraft on the Sun. They found that such distributions are slightly offset to the west from the solar source. Long-lived particle injection from CME-driven shocks has been suggested as the reason for this shift (Lario et al. 2014). They also estimated that the shock height is within approximately 40 Rs when the particle injection from the shock maximizes. Given the observation that CME-driven shocks form at a heliocentric distance of 1.5 to 5 Rs (Tsurutani et al. 2003a; Gopalswamy et al. 2013b), these findings indicate that observations close to the Sun are most important for large SEP events. Furthermore, the range of heliocentric distances matches with the frequency range of IP type II bursts (Gopalswamy et al. 2012c).
SEP intensity variability
Attempts to understand the variability in SEP intensities have been focused on two aspects: (1) source factors involving CME properties, (2) environmental factors such as ambient magnetosonic speed, preceding CMEs, deflection by large-scale structures, seed particles, and turbulence. The CME speed and width are the primary source parameters. In cycle 23, about approximately 75% of SEP-producing CMEs were found to be halos (Gopalswamy 2006b) for the whole cycle as well as for the first 62 months. On the other hand, 100% of SEP-producing CMEs are halos in cycle 24 so far. This remarkable result can be understood in terms of the anomalous expansion of CMEs in cycle 24 due to the reduced heliospheric pressure. The average sky-plane speed of SEP-producing CMEs is approximately 1,500 km/s, slightly higher than that in cycle 23 over the corresponding phase of the solar cycle (Gopalswamy 2012). These results lead to the conclusion that the cycle-24 CMEs need to have more kinetic energy to produce SEP events similar to those in cycle 23. In other words, the cycle-24 CMEs seem to be less efficient in accelerating particles.
SEP events associated with weak and strong eruptions
Solar cycle effects
It should be noted that the largest SEP events are associated with quasi-parallel shocks (Kennel et al. 1984a, 1984b). This is because of the presence of upstream turbulence ahead of the shocks (Tsurutani et al. 1983), leading to enhanced Fermi-type acceleration across the shocks.
The paucity of GLE events in cycle 24
Ground-level enhancement (GLE) in SEPs represent the highest energy particles accelerated by the Sun, making the particles penetrate Earth’s atmosphere to the troposphere where they produce air showers like those produced by the galactic cosmic rays (GCRs) and the air showers are detected by ionization chambers, muon detectors, and neutron monitors on the ground. These hard-spectrum events affect navigation systems, spacecraft electronics and operations, space power systems, manned space missions, and commercial aircraft operations (Shea and Smart 2012; Mewaldt et al. 2012; Kudela 2013; Ruffolo et al. 2013). In particular, they can be a significant radiation exposure to humans in space and in airplanes on polar routes. GLE events also have important implications for VLF wave propagation: Zigman et al. (2014) found that at energies up to approximately 2 GeV, the ionization rate for solar protons may exceed the GCR ionization by 1.5 orders of magnitude.
Typically about a dozen GLEs occur in each solar cycle, which is about 15% of the number of large SEP events during cycles 19 to 23 (Shea and Smart 2008). However, there were only two GLEs during cycle 24 so far, even though there were 34 large SEP events, amounting to <6% (Gopalswamy et al. 2013a; Thakur et al. 2014; Papaioannou et al. 2014). The paucity of GLE events in cycle 24 cannot be explained by the 22% reduction in the number of fast and wide CMEs originating in the traditional GLE longitudes (W20 to W90). A combination of several factors is needed to explain the drastic reduction in the number of GLE events: (1) The reduction in the ambient magnetic field in the corona reduces the efficiency of shock acceleration especially for GeV particles, even though it is not a problem in accelerating approximately 10-MeV particles. (2) Many CMEs have nonradial motion either due to the coronal environment or inherent asymmetry in the source region. The nonradial motion makes the shock nose not well connected to Earth (Gopalswamy et al. 2013a; Gopalswamy and Mäkelä 2014). This means that even if GeV particles were accelerated at the shock nose, they may not reach the observer. (3) The CME may be ejected into a locally tenuous corona making the shock weaker.
The fact that both CMEs in Figure 29 were large SEP events raises an important question on the location of particle acceleration on the shock surface. The stringent requirement of the ecliptic distance for GLE CMEs suggests that GLEs may be accelerated at the shock nose, where the shock is the strongest. On the other hand, magnetic connection to any part of the shock seems to suffice for large SEP events. This means that lower energy SEPs are accelerated over most of the shock surface. This is also consistent with the study by Dalla and Agueda (2010) who found that the probability of detecting SEP events remained constant up to a latitude of 28°. Gopalswamy and Mäkelä (2014) examined the latitudinal connectivity issue in historical GLE source regions that were at latitudes >30°. They were able to show that in all the higher latitude cases, there was a polar coronal hole, a streamer, or a pseudo-streamer poleward of the GLE source region suggesting deflection toward equator and hence enhancing the possibility of shock nose connection to Earth.
The requirement that the CME nose be in the ecliptic for GLE events may also have implications for the shock geometry. It is reasonably certain that GLEs are released when the CME is at a height of approximately 3 Rs. Since the shock nose has the largest heliocentric distance compared to other parts, it is likely that the nose is above the source surface and the upstream field magnetic field is open. This suggests that GLE-producing regions of the shock may have quasi-parallel geometry. Recent investigations confirm the suggestion that CME-driven shocks are likely to be supercritical and quasi-parallel near the nose, while subcritical and quasi-perpendicular at the flanks (Bemporad et al. 2014; Bemporad and Mancuso 2011, 2013). It is well known that the first critical Mach number is the smallest (approximately 1.5) for quasi-parallel shocks and the largest for quasi-perpendicular shocks (approximately 2.7) (see e.g., Mann et al. 2003). Supercritical shocks are known to accelerate particles more efficiently (Burgess and Scholer 2013). Thus for a given shock speed (approximately 2,000 km/s for GLE events) and the typical Alfvén speed of approximately 600 km/s at 3 Rs (Gopalswamy et al. 2001b), the quasi-parallel shocks are likely to be definitely supercritical, while the quasi-perpendicular shocks may or may not be supercritical. These conclusions on the shock geometry and connectivity need to be further explored and modeled for a full understanding of particle acceleration by CME-driven shocks.
Extreme space weather
Given that CMEs are responsible for the largest geomagnetic storms and SEP events, it is natural to think that some properties of CMEs or their source regions may result in extreme events. If we define extreme events as those, which lie on the tail of a distribution, we can readily identify the CME’s speed as one critical parameter. From the cumulative distribution of CME speeds, one infers that there are not many CMEs with speeds exceeding approximately 3,500 km/s (e.g., Gopalswamy et al. 2010b). When a CME has a speed >3,500 km/s, it may be thought of as an extreme event. The consequences of such an energetic CME are likely to be extreme also. An extremely fast CME will drive a shock, which will accelerate particles to very high energies. If all conditions for the acceleration of particles are conducive (low ambient magnetosonic speeds, good connection of the shock nose to the observer, a quasi-parallel shock, and a preconditioned ambient medium), one might expect an extreme SEP event. Kovaltsov and Usoskin (2014) determined the cumulative occurrence probability distribution of SEP events based on directly measured SEP fluences for the past 60 years, estimates based on the terrestrial cosmogenic radionuclides 10Be and 14C for Holocene time scale, and cosmogenic radionuclides measured in lunar rocks on a time scale of up to 1 Myr. They concluded that SEP events with a >30-MeV proton fluence greater than 1011 (protons cm−2 year−1) are not expected on a Myr time scale.
The SEP event depends just on the outer structure of a fast CME, viz., the shock. It should be noted that the most intense interplanetary shock detected at 1 AU distance from the sun was detected on March 8 in the 7 to 17 March 2012 CAWSES-II study interval and had a magnetosonic Mach number of approximately 9.4 (Tsurutani et al. 2014). However, the maximum shock Mach number is possible, assuming an ICME speed of 2,700 km/s is 45 (Tsurutani and Lakhina 2014). One speculation to explain the difference is that it is possible that the acceleration of energetic particles at the shock leads to ‘shock damping,’ reducing the shock intensity. Theoretical studies are needed to confirm/deny this hypothesis.
On the other hand, the generation of a large geomagnetic storm depends on the internal structure of the CME, especially the magnetic cloud (MC, Burlaga et al. 1982) and filament (Kozyra et al. 2014) internal to it. What is important is that the magnetic field intensities are high and southward. It has been shown empirically that fast CMEs have particularly intense magnetic fields (Gonzalez et al. 1998). Sheath magnetic fields (Tsurutani et al. 1988) are swept up slow solar wind fields and not part of the CME proper. Sheath fields can have intensities of only approximately four times the upstream ambient field strength (Kennel et al. 1985; Tsurutani and Lakhina 2014) unless multiple shocks are compressing the sheath, so this region is, in general, unimportant for the cause of extreme storms. If the magnetic cloud has its axis pointing to the south and the magnetic field is high, one could expect strong magnetic interconnection between the interplanetary magnetic field and the Earth’s magnetopause magnetic field (Dungey 1961) and therefore an intense geomagnetic storm. In addition, if the shock remained strong, it would produce an extreme impulse on the magnetosphere that can expose geosynchronous orbit to the solar wind (Tsurutani and Lakhina 2014) and a sudden (positive) impulse (SI+) on the ground (Tsurutani et al. 2008). An intense shock would also produce an extreme energetic storm particle (ESP) event at Earth (Tsurutani et al. 2009b). Thus for solar terrestrial relationships, an extreme event in its origin will also have extreme consequences at Earth and near-Earth space environment.
Tsurutani and Lakhina (2014) considered an extreme scenario for geomagnetic response. Since the CME speed and its magnetic content can ultimately be traced to the solar source (Gopalswamy 2010), an extreme CME would be born in an active region with enormous free energy. Going by the highest field strength ever observed in a sunspot (approximately 6,100 G, Livingston et al. 2006) and the largest active region area (5,000 msh, Newton 1955), one can estimate a free energy of approximately 1036 erg. This is two to three orders of magnitude larger than the amount of free energy estimated for NOAA AR 10486, which produced several of the Halloween CMEs (Gopalswamy et al. 2005c). If the entire approximately 1036 erg goes into the kinetic energy of a CME of mass approximately 1018 g, then the CME speed would be approximately 14,000 km/s, about four times larger than the CME speed ever observed.
Recently, Maehara et al. (2012) investigated 365 stellar flares with energies in the range 1033 to 1035 erg using the Kepler mission data. Fourteen of those flares were from Sun-like stars (G-type main sequence stars with rotational periods >10 days and surface temperature in the range 5,600 to 6000 K). They estimated that superflares with energy 1035 erg occur once in 5,000 years. Note that this energy is an order or magnitude less than the maximum free energy estimated by Gopalswamy (2010). It must be pointed out that CMEs are not observed directly in stellar eruptions, but if the eruption mechanism is similar to that on the Sun, one would expect even higher CME kinetic energies than the 1035 erg flare energy found by Maehara et al. (2012). In fact, the CME kinetic energy gains the largest share among various ways in which the free energy in an active region is divided (Emslie et al. 2012). However, one should take note that the amount of energy that goes into the flare particles is deducted from the CME energy (Tsurutani and Lakhina 2014).
Shibata et al. (2013) also investigated theoretically the possibility of superflares on the Sun using the current dynamo model and concluded that it is indeed possible to generate magnetic flux necessary for producing superflares of energy in the range 1034 to 1035 erg within the next 40 years. The Solar Evolution and Extrema (SEE) project of the VarSITI (Variability of the Sun and Its Terrestrial Impact) program of SCOSTEP is expected to make further progress on this issue over the next few years.
The Carrington flare of 1859 remains one of the important benchmark events for extreme solar events. The flare itself was detected with naked eyes, indicating that it is certainly an unexpected event. Cliver and Dietrich (2013) estimated that the soft X-ray flare size of X45, compared to X35 for the cycle-23 flare on 4 November 2003. The Carrington eruption did produce a historical geomagnetic storm, the highest intensity on record: Dst = −1,760 nT (Tsurutani et al. 2003b). Kataoka (2013) estimated that the probability of a Carrington-type geomagnetic storm occurring within the next decade is approximately 4% to 6%. Yermolaev et al. (2013) performed a statistical analysis of the OMNI data for the period 1976 to 2000 and concluded that a Carrington-type event could occur once every 500 years (see also Alves et al. 2011). From the flare onset to the geomagnetic storm onset, Carrington (1859) gave a time of approximately 17.5 h from the flare to the storm, which indicates a CME speed of approximately 2,360 km/s (Gopalswamy et al. 2005c). The all-time record for the shortest CME transit time from the Sun to 1 AU was the 12 August 1972 event (Vaisberg and Zastenker 1976) which was approximately 14.6 h. With such a fast transit time and large magnetic cloud magnetic field, why did not a large magnetic storm occur? It was discovered that the magnetic field orientation within the magnetic cloud was northward, so the Earth had geomagnetic quiet rather than a large storm (Tsurutani et al. 1992a). It has been shown that for northward IMFs impinging on the Earth’s magnetosphere, one has an extreme geomagnetic quiet rather than an intense storm (Tsurutani et al. 1992b; Du et al. 2011). These CME transit times are relatively close to the maximum time of 12 h, assuming an initial CME speed of 3,000 km/s (Tsurutani and Lakhina 2014). Two of the Halloween storms in cycle 23 had such short transit times: 18.9 and 19.7 h (Tsurutani et al. 2005; Mannucci et al. 2005) for the 28 October 2003 and 29 CMEs, respectively. In the extremely weak cycle 24, there was another extreme event on 23 July 2012 observed by multiple spacecrafts. This well-observed event merits some additional discussion.
23 July 2012 event
On 23 July 2012, the GOES satellite reported a large SEP event with a peak >10-MeV flux of approximately 12 pfu (1 pfu = 1 particle sr−1 cm−2 s−1), a very ordinary event. However, the solar source of the associated CME was approximately 45° behind the west limb and the CME was heading roughly toward STEREO-A (located at W121), which observed the CME as a full halo. STEREO-A also detected a 5,000-pfu SEP event with an ESP event approximately ten times larger, similar to the 23 March 1992 and 20 October 1989 events (43,000 and 40,000 pfu, respectively). The CME arrived at STEREO-A in approximately 19 h, making it one of the historical events potentially greater than the Halloween 2003 events (Russell et al. 2013b; Baker et al. 2013; Mewaldt et al. 2013; Gopalswamy et al. 2014b). The STEREO-A SEP event started approximately 4 h earlier than the GOES event and approximately 8 h earlier than the STEREO-B event because STEREO-A was better connected to the source than the other two spacecraft. There are a number of similarities between the 23 July 2012 event and the 20 October 1989 event studied by Lario and Decker (2002).
If the CME were heading toward Earth instead of STEREO-A, it would have caused another geomagnetic storm of historical proportions (Liu et al. 2014b). The total magnetic field strength was approximately 109 nT in the shock sheath and approximately 60 nT in the ICME. The southward component had a peak value of −52 nT. Liu et al. (2014) estimated that the Dst index to be in the range −1,150 nT to −600 nT. The simple empirical relation reported by Gopalswamy (2010) also gives a high Dst index: Dst = −0.01 VBz - 32 nT. Plugging in V=1,500 km/s and Bz = 52 nT, the Dst index becomes −812 nT. This is half of the strength reported for the Carrington storm (−1,760 nT, Tsurutani et al. 2003b) and a quarter of the maximum possible of Dst = −3,500 nT (Tsurutani and Lakhina 2014). That is if the magnetosphere does not saturate at Dst = −2,500 nT as Vasyliunas (2011) has argued. If we compare with the estimate of Siscoe et al. (2006), who used hourly averages to estimate the Dst index as −850 nT, we see that the July 2012 storm has the same strength as the Carrington event (see also Cliver and Dietrich 2013 who estimate Dst approximately −900 nT). The important point is that such extreme storms can occur even during the subdued heliospheric conditions prevalent in cycle 24. It is possible that the July 2012 storm already represents a reduction in Bz because of the anomalous CME expansion; otherwise the storm could have been even bigger than the current estimates. If the ICME were south pointing with the 109-nT field during the July 2012 event (Liu et al. 2014) and the speed were similar to the shock speed (2,250 km/s), the empirical formula (Dst = 0.01 VBz - 32 nT) would predict a storm of approximately −2,500 nT.
1859 Carrington ionospheric event
In panel a, prior to the event, the two EIAs are clearly noted one spanning approximately −5° to −30° and a second from approximately +5° to 20° LAT (we use a 3.25 × 106 ions/cm−3 to define the enhancement areas). Panel c shows that the oxygen peaks are located at approximately 530 to 920 km and approximately 500 to 900 km for the southern and northern regions, respectively. The magnetic latitude ranges are approximately −20° to −35° LAT and +30° to +45° LAT. The peak values reach approximately 6 × 106 O+ cm−3 at altitudes centered at approximately 700 km. It is noted that the densities of the uplifted EIA peaks at 850 km and 1,000 km were approximately 4 × 106 and approximately 3.5 × 106 cm−3, respectively. These latter ion densities are substantially greater than quiet-time neutral density values, being approximately 40 times at 850 km and approximately 300 times at 1,000 km. Thus, low altitude satellite drag will be substantially increased in such an extreme storm.
The weak activity as the Sun climbed toward the maximum phase of solar cycle 24 has provided an enormous opportunity to study the Sun and the heliosphere under quieter conditions. The space weather is less severe compared to cycle 23 in terms of intense geomagnetic storms and powerful SEP events. On the other hand, the cosmic-ray intensity remains exceptionally high after reaching the highest levels in the space age during the extended solar minimum. The geomagnetic ap index was the lowest on record. The latter was not only due to the changes of the Sun but also the location of the Earth relative to solar active regions (coronal holes). The vast array of ground- and space-based instruments has helped define this altered state of the Sun and heliosphere. The CAWSES-II program enabled many researchers to take the system approach and be mindful of the implications of their research for solar terrestrial relationship. From the eruptive events during the weak solar cycle to the dynamo problem and the grand minima of solar activity, the Sun has triggered a plethora of research activities, including the MiniMax24 activity that recognizes the importance of studying weak activity. The VarSITI program will address the declining phase of the weak cycle 24, when there will be phase-specific phenomena such as frequent CIRs and high-speed streams from the Sun.
The TG3 activities of the CAWSES-II program will naturally connect to the International Studies of Earth-affecting Transients (ISEST) project of VarSITI. Campaigns are being conducted in observing transient solar events under the MiniMax24 component of the ISEST project. The working groups under ISEST will deal with the theory, modeling, and detailed analysis of the campaign events. There will also be a close connection to the Solar Evolution and Extreme (SEE) project in examining flares on sun-like stars and extreme events such as the Carrington event and the 23 July 2012 event. These connections provide continuity to the SCOSTEP science activities and make progress without serious interruptions.
atmospheric imaging assembly
climate and weather of the Sun-Earth system II phase
CME analysis tool
coordinated data analysis workshop
corotating interaction region
coronal mass ejection
field of view
galactic cosmic rays
ground level enhancement
Geostationary Operational Environmental Satellite
high-intensity long-duration continuous geomagnetic activity
interplanetary coronal mass ejection
large angle and spectrometric coronagraph
Scientific Committee on Solar Terrestrial Physics
Solar Dynamics Observatory
solar energetic particles
solar geophysical data
Solar and Heliospheric Observatory
Solar Terrestrial Relations Observatory
total electron content
Task Group 3 (of CAWSES-II)
total solar irradiance
The authors thank Kazunari Shibata and Joseph Borovsky for providing leadership to the Task Group 3: How does short-term solar variability affect the geospace environment? The authors also thank CAWSES-II leaders Susan Avery, Alan Rodger, Toshitaka Tsuda, and Joseph Davila for their effort in running a successful SCOSTEP scientific program. We thank P. Mäkelä and Jing Huang for carefully reading the manuscript. We also thank Kazunari Shibata for his comments on the manuscript. The work of NG was supported by NASA’s LWS TR&T program. Portions of this research were performed at the Jet Propulsion Laboratory, California Institute of Technology under contract with NASA. We thank the anonymous referees for the helpful comments.
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