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ISSN 0001 4338, Izvestiya, Atmospheric and Oceanic Physics, 2010, Vol. 46, No. 7, pp. 799­819. © Pleiades Publishing, Ltd., 2010. Original Russian Text © Yu.I. Yermolaev, M.Yu. Yermolaev, 2010, published in Geofizicheskie protsessy i biosfera, 2009, Vol. 8, No. 1, pp. 5­35.

Solar and Interplanetary Sources of Geomagnetic Storms: Space Weather Aspects
Yu. I. Yermolaev and M. Yu. Yermolaev
Space Research Institute, Russian Academy of Sciences, Profsoyuznaya ul. 84/32, Moscow, 117997 Russia Abstract--The current notions of the solar­terrestrial relations responsible for the transport of solar distur bances and for the generation of magnetic storms on the Earth are briefly reviewed. The probability of gener ating magnetic storms by different solar and interplanetary phenomena is quantitatively estimated. The effi ciencies of generating magnetic storms by different types of solar wind streams are compared. Keywords: magnetic storms, flares, coronal mass ejections, magnetic clouds. DOI: 10.1134/S0001433810070017

INTRODUCTION With advances in technical potentialities and with our extended knowledge of nature, it is becoming more and more evident that space factors affect not only different space and ground based technological systems [Lilensten, 2007; Plazmennaya Geliogeofizika (Plasma Heliogeophysics), 2008], but they also signifi cantly affect biological objects, including the human organism [Plazmennaya Geliogeofizika (Plasma Helio geophysics), 2008]. When studying the influence of space factors, the term "space weather" often implies a combination of phenomena that are physically rather heterogeneous; therefore, in every particular case, the use of this term needs a more precise definition. In fact, a new scientific field--solar­terrestrial rela tions--has been formed. In the framework of this direction, all possible interactions between helio and geo physical phenomena are studied. Figure 1 sche matically shows the structure of solar­terrestrial rela tions. It is almost impossible to find a systematic descrip tion of the basic principles of this scientific field in any domestic or foreign publication, because it is both inter and multidisciplinary and includes the elements of a number of sciences. These elements are most often presented in the specialized literature on one or another of scientific field and often fall through the cracks in regards to the specialists of related directions. A more detailed description of the problems related to the solar­terrestrial relations can be found in the above mentioned monograph [Plazmennaya Gelio geofizika (Plasma Heliogeophysics), 2008] and in the encyclopedia edited by R.A. Syunyaev [Fizika Kos mosa (Space Physics), 1986], which is still of actual although it was published 20 years ago. It is important to note that both books are published in Russian, which makes them accessible to Russian researchers.

We think that, before proceeding with our basic results, we should give a brief popular description of some basic elements of the system under consider ation. In our opinion, such a description will give even laymen a better insight into what follows. In addition, there is no commonly accepted terminology, which presents some difficulties in discussing the problems of the solar­terrestrial relations. To denote the described processes and phenomena, we use abbreviations taken from the scientific literature in English. 1. BASIC TERMS AND DEFINITIONS When speaking about solar­terrestrial relations, it is necessary to emphasize that there are two channels of energy transfer from the Sun to the Earth: electro magnetic and corpuscular radiations. The former is considered basic: it is through this channel that most solar energy is transferred to the Earth (about 1.37 kW per every square meter of the land surface). This energy flux lies mainly within both visible and infrared wavebands and is characterized by its steadiness; its variations do not exceed fractions of percent, and, therefore, it is referred to as a solar constant. Reaching the Earth in over 8 min, this flux, which is absorbed mainly by the atmosphere and the land surface, plays an important role in atmospheric weather. The electromagnetic radiation within both X ray and ultraviolet bands significantly varies during the development of active processes occurring on the Sun. The energy fluxes within the indicated bands are extremely small: even when, during the strongest solar flares, the X radiation flux increases by three orders of magnitude, the total energy flux remains six orders of magnitude smaller than the solar constant. In this case, it should be remembered that the indicated radi ations are almost completely absorbed by the Earth.

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The Sun

Electromagnetic radiation: visible light, ultraviolet radiation, radio radiation, X ray radiation, and others

Corpuscular radiation: the solar wind and solar cosmic rays

Interplanetary medium Neutral atmosphere Magnetosphere and ionosphere Biological effects Geomagnetic field

Meteorological effects

Processes inside the Earth

Fig. 1. Structure of solar­terrestrial relations.

The latter channel--corpuscular radiation--is dominant in "space weather", and it is precisely this channel which will be considered below. Corpuscular radiation consists of solar wind and cosmic rays. In recent years it has been common prac tice to call cosmic rays energetic particles. This name better reflects their physical nature, because cosmic rays are charged particles (electrons, protons, and ions) accelerated to very high (often relativistic) veloc ities. These particles are of galactic or solar origin. The former particles are born outside the solar system. On average, their occurrence on the Earth's orbit is less frequent than the occurrence of particles of solar ori gin. An increase in the Sun's activity results in a decrease in the flux of these particles. During active processes occurring on the Sun (flares, destruction of arcs, coronal ejections, etc.) and in the interplanetary medium (mainly on shock waves), energetic particles of solar origin accelerate. Basically, accelerated particles are radiation that can penetrate into bodies and destroy the molecules of both animate and inanimate natures. Fortunately, the Earth's surface is safely protected by the magneto sphere and the atmosphere. However, during space and even airplane transarctic flights, energetic parti cles may be hazardous to people and electronic devices. It is under the influence of radiation that most instruments installed aboard spacecrafts fail to oper ate. For example, in October­November 2003, some malfunctions of the instrumentation installed aboard the SOHO and ACE spacecrafts were associated with this reason [Veselovsky et al., 2004].

The temperature of the solar corona's plasma amounts to 2 â 106 K, and, as a result, this plasma can not be completely confined by the Sun's gravitational field, "escapes" to the interplanetary space, and fills the heliosphere with itself. Although almost the entire solar system is within the solar corona, plasma that is more than a few solar radii away from the Sun and that has characteristics significantly different from those of the plasma at the corona's base is usually called solar wind. Having a mean velocity of 400 km/s, the solar wind reaches the Earth over 2­5 days. In this case, its density in the Earth's orbit amounts to a few ions per 1cm3, which is impracticable under the conditions of ground based experimental installations. Neverthe less, solar wind has a profound effect on the energy transfer from the Sun to the Earth's magnetosphere and its outer layers. Slow changes that occur in the system under con sideration and that are characterized by the time on an order of months and more are sometimes called "space climate." If they are eliminated from consider ation, the dynamic portion remains, which is charac terized by rapid deviations from an averaged pattern that is the subject of investigation in studying space weather. We will restrict ourselves to a description of only a small part of the given scheme--the transfer of distur bance from the Sun to the Earth's magnetosphere through the solar wind. In this case, most attention will be concentrated on recent results obtained from studies of the sources of the strongest magnetospheric disturbances (magnetic storms on Earth).
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In recent years, in heliobiophysics that concerns the problems associated with the influence of space weather factors on biological objects, including human beings, studies of the role of magnetic storms have taken on greater importance. In this direction, significant results have already been obtained [Vil loresi et al., 1994; Gurfinkel', 2004; Zenchenko and Breus, 2008; and Kleimenova et al., 2008] and ways have been outlined to decrease the risk of serious con sequences for people through taking preventive mea sures with the approach of magnetic storms. However, the number of similar investigations carried out thus far is insufficient, mainly because they are complex and interdisciplinary and require definite knowledge in the related scientific areas. In particular, these investigations were limited to a comparison with the very fact of the occurrence of magnetic storms. Cur rently, there are no serious comparisons of biological responses to magnetic storms with their properties and origin, because such comparisons require knowledge of magnetic storms and their sources. Heliobiophysi cists are, as a rule, not very aware of the Sun's physics and solar­terrestrial relations. Their use of indices of geomagnetic activity and their classification according to activity level are formal and not always physically justified. This paper intended for a wide circle of read ers has been written by us to clarify a number of ques tions related to the origin and character of geomag netic storms, classification of the geomagnetic activ ity indices, and to the adequate use of these indices in solving different applied problems. From the preceding, it follows that the study of solar and interplanetary sources of geomagnetic storms remains an relevant and important problem of space weather and its numerous applications. Our general notion of the sources of storms has not changed over many years: the basic source of mag netospheric disturbances is the negative (southward) Bz component of the interplanetary magnetic field (IMF), because, in this case, the magnetosphere becomes open and energy from the solar wind can arrive in the magnetosphere and result in magnetic storms. The IMF usually lies within the ecliptic plane and does not contain any of the Bz components; only the disturbed types of the solar wind can contain the IMF Bz component, including the southward one. According to current views, there are two basic chains (scenarios) of energy transfer from the Sun to the magnetosphere. (Since the terminological diffi culties of the above described scientific field have already been noted, below we will give the English names of each process and phenomena and use the abbreviations derived from their English names). Sce nario 1: solar disturbance (solar flare and coronal mass ejection (CME)) interplanetary CME (ICME, ejecta, and magnetic cloud (MC)), including the IMF southward Bz component, magnetic storm. Sce nario 2: coronal holes that form high speed solar wind streams interaction of a high speed flow with the
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preceding low speed flow and the formation of the IMF compression and deformation region (corotating interaction region (CIR)), which includes the IMF southward Bz component magnetic storm. Although the mechanisms of energy transfer have been studied over many years and, by now, a significant vol ume of both experimental and theoretical data has been accumulated (see, for example, [Lilensten, 2007; Schwenn, 2006; Pulkkinen, 2007] and their refer ences), this problem has yet to be finally solved. On the one hand, the investigations that were car ried out include a long chain of spatial regions with different physical processes, they are of interdiscipli nary character, and they require the joint efforts of sci entists of different specialities. On the other hand, among the regions known to us, there are some zones on which we have no experimental data and we can only propose hypotheses for their interrelations. For example, there are data on the Sun and circumsolar space obtained with the remote sensing methods and there are data obtained from direct measurements in the near Earth space; however, there are almost no data on the region between the circumsolar and near Earth spaces for lack of measurements in this region. We also know almost nothing about the thin fronts of a bow shock and the magnetopause because the motion of these boundaries is too fast with respect to the satel lite. Therefore, in this paper, we will not consider regions for which experimental data are available; they are discussed in detail in the special literature. We will focus our attention on the "interface" between these regions, which is usually absent in the special litera ture. In our previous papers [Yermolaev and Yermolaev, 2003b, 2006; Yermolaev et al., 2005a], it was shown that the quantitative relations between different phe nomena depend strongly on the methodical approaches used. Therefore, problems related to the influence of the quantitative determinations of phe nomena and ways to compare them on the estimates of correlation between these phenomena will be dis cussed in the following sections. Then, some estimates of correlations established on the basis of a large body of observational data will be given. Finally, it will be shown that, in most cases, the generation of magnetic storms is characterized not only by the IMF southward Bz component but also by a certain behavior of other solar wind parameters. This allows us to suggest that the magnetic storms induced by a disturbed plasma compression region before ejecta/MC (sheath), MC, and CIR can be generated through different physical mechanisms. It should be noted that there is a double meaning of the word geoeffectiveness. In one case, geoeffectiveness implies a probability with which one or another phe nomenon can cause a magnetic storm, i.e., the ratio between the number of events of a chosen type result ing in a magnetic storm and the total number of these events. In the other case, geoeffectiveness implies the
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3

2

1 0 M0

M5

X0

X5

X10

X15

X20

Fig. 2. Ratio between the optical (vertical axis) and X ray (horizontal axis) values (classes) of solar flares. 643 flares of class M5 and higher observed in 1976­2000 have been analyzed [Yermolaev and Yermolaev, 2003b, 2006]. The dashed line denotes the linear approximation of the given data.

efficiency of storm generation by unambiguously interrelated phenomena, i.e., the ratio between the "output" and "input" of a physical process, for exam ple, between the values of the Dst index and the IMF southward Bz component. 2. PHENOMENA ON THE SUN Data on phenomena occurring on the Sun are of specific character. Unlike the data on interplanetary and magnetospheric phenomena obtained from in situ measurements of their parameters, data on solar phe nomena are obtained from remote (ground based or near Earth space based) measurements of the solar atmosphere within different frequency ranges of elec tromagnetic waves. The frequency of radiation is determined by conditions in the radiating volume of plasma (mainly, by its concentration), and, generally speaking, the measurements taken in different fre quency ranges yield the characteristics of the Sun's different regions. Determining the dynamics of a solar phenomenon, including its spatial motion (especially along the line of sight) is a complicated and ambiguous problem, because, in this case, different phenomenon components whose characteristics and locations vary in time must be measured by different instruments. In this case, it is assumed that the results of measure ments with different instruments can be used in study ing one and the same phenomenon. Solar flares were first measured within the optical wave band with ground based instruments, and it is on the basis of optical observations that solar flares were classified (see, for example, [Krajcovic and Krivsky, 1982]). However, the orbital observations of the Sun within the X ray band (which are impossible for ground based measurements) were made possible with the advent of the space age; the X ray flares were classified on the

basis of (GOES) satellite measurements (see http:// www.ngdc.noaa.gov/stp/GOES/goes.html). The opti cal and X ray emissions are formed at different stages of solar flares and in their different regions. Thus, the flare values (classes) determined with two different methods are of different physical natures. The ratios between the optical and X ray values of solar flares are given in Fig. 2 for the period 1976­ 2000. All flares with X ray values of M5 and higher, which are usually treated as candidates for the sources of interplanetary and magnetospheric disturbances, are shown [Yermolaev and Yermolaev, 2003b, 2006]. Figure 2 clearly shows that there is correlation only in a statistical sense, because some events can simulta neously have a high optical class and a low X ray class and vice versa. Over a long period of time, all disturbances within the solar wind and the Earth's magnetosphere were associated only with solar flares. Figure 3 (on the left) shows all X ray flares of both high (M) and extreme (X) classes (gray and black squares, respectively). To make the flares, whose lengths were from a few minutes to a few dozen minutes, distinguishable in Fig. 3, their durations were increased up to 6 h in plotting the graph. Figure 3 (on the right) shows both average (­50 > Dst > ­100 nT) and strong (Dst < ­100 nT) magnetic storms (gray and black squares, respectively); their durations in the graph correspond to those observed in reality. The numbers of the days of the sun's Carrington rotations (about 27 days) are plotted along the abscissa axis, and the years from 1976 to 2000 are plotted as ordinates. On the whole, on the time scales of the Sun's several rotations, a good correlation is observed between solar and magnetospheric events. However, in most cases, the attempts to relate concrete events to one another prove to be ineffective; this will be discussed in more detail in Section 4. In the early 1970s, one more powerful solar phe nomenon--CME--was revealed with white light coronagraphs installed aboard spacecrafts. Over a long period of time, the CMEs were studied only by indi vidual researchers and were almost neglected in con sidering the chain of the solar­terrestrial relations. However, after the publication of Gosling's paper in 1993, the situation changed, and now the CMEs are treated as the single cause of all interplanetary and geomagnetic disturbances, although both the physical phenomena (flares and CMEs) are closely interrelated (see, for example, discussions in [Harrison, 1996; Cliver and Hudson, 2002; and Yashiro et al., 2005]. A large body of CME data was obtained with a LASCO coronagraph on board the SOHO spacecraft (http://cdaw.gsfc.nasa.gov/CME_list/). When study ing the geoeffectiveness of CMEs (unlike flares that can be observed on the solar disc), a very important problem is to determine the location of their source on the solar disc and, first of all, answer the question of what side of the solar disc their source is located on visible or back side. CMEs are observed in two dimen
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SOLAR AND INTERPLANETARY SOURCES OF GEOMAGNETIC STORMS 0 2001 2000 1999 1998 1997 1996 1995 1994 1993 1992 1991 1990 1989 Carrington rotation 1988 Year 1987 1986 1985 1984 1983 1982 1981 1980 1979 1978 1977 1976 0 5 10 15 20 25 Day of the Sun turn Flares M Flares X 1650 1700 1750 1800 1850 1900 1950 5 10 15 20 25 2001 2000 1999 1998 1997 1996 1995 1994 1993 1992 1991 1990 1989 1988 Year 1987 1986 1985 1984 1983 1982 1981 1980 1979 1978 1977 1976 0 10 15 20 25 Day of the Sun turn ­100 < Dst < ­50 Dst < ­100 5 1650 1700 1750 1800 Carrington rotation 1850 1900 1950 0 5 10 15 20 25

803

Fig. 3. Time variations in solar flares (on the left) and magnetic storms (on the right) for the period 1976­2000.

sional images in which the Sun is "cut out" of the viewing field of coronograph due to its occulting disc. To solve this problem, the results of observations of CMEs in the white light outside the solar disc are com pared with the results of observations of other solar phenomena, such as flares, ultraviolet and X ray dim mings, ultraviolet luminosities, both ultraviolet and X ray posteruptive arcades, etc., on the solar disc in
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different wavebands. Such a comparison between the results of observations in the white light and ultraviolet band is shown in Fig. 4 [Gopalswamy, 2002]. Thus, it should be remembered that the CME location deter mined with the above indicated method is not an experimental fact but a hypothesis, because, for this purpose, researchers have to use the results of mea surements taken (i) with different instruments; (ii) on
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()

(b)

Fig. 4. Superposition of the images of CME moving toward the Earth on July 14, 2000 (the Bastille Day event) in the white light (SOHO/LASCO, SOHO/EIT). The background in the form of "white snow" (on the right) results from a bombardment of the SOHO detectors by energetic particles associated with solar activity [Gopalswamy, 2002].

different frequencies; (iii) in different spatial regions; and, finally, (iv) at different times. Therefore, the loca tion of the CME source on the solar disc can only sta tistically be considered on the basis of images obtained within other wavebands. There are experimental facts that some CMEs lead to direct measurements of ejecta/MC (interplanetary CME) and magneto spheric disturbances, but they do not have any appar ent features on the Sun's visible disc [Zhang et al., 2003]. If such CMEs are neglected, then, on the basis of only solar observations, they can be included in the list of CMEs on the invisible side of the Sun and can lead to wrong conclusions about the CME geoeffec tiveness [Yermolaev, 2008]. Unlike flares and CMEs, coronal holes are suffi ciently stable solar structures that can exist for several 27 day solar rotations. Coronal holes have an open configuration of the magnetic field which allows the corona to form fast solar wind streams (Fig. 5).
Table 1. Quasistationary types Type 1 Heliospheric current sheet Type 2 Slow wind from coronal streamers Type 3 Fast wind from coronal holes Disturbed types Type 4 Compressed solar wind streams (compression region between slow and fast flows (CIR) and compression region before MC and ejecta (sheath)) Type 5 Interplanetary CME (MC and ejecta) Type 6 Rarefaction region

3. INTERPLANETARY PHENOMENA The classification of large scale events in the inter planetary medium arose with the advent of the space age and is developing now as knowledge of and data on the solar wind and its sources on the Sun accumulate. Although the classification methods are developing rapidly enough now, a general idea of the types of solar wind has not changed significantly. According to a large body of observational data, there are six basic types of large scale solar wind streams (Fig. 6, Table 1). Among the types given in Table 1, only two types 4 and 5 are geoeffective, because they can include the long southward IMF Bz component [Gosling and Pizzo, 1999; Gonzalez et al., 1999; Crooker, 2000; and Bothmer, 2004]. An analysis of literature shows that there is no one method for identifying interplanetary phenomena; different researchers use different sets of parameters and different numerical criteria in analyzing these phenomena. For example, to identify a magnetic cloud, the available methods may include from two to ten parameters (see, for example, [Yermolaev and Yer molaev, 2003b] and the references from this paper). In the literature, there are several lists of ICMEs (MC and ejecta) [Cane and Richardson, 2003; Zhang et al., 2004; and Echer et al., 2005] and one list of CIRs [Alves et al., 2006], but there are no lists of other types of solar wind streams or lists that simultaneously include different types of streams. We have prepared a list of all the above indicated types of solar wind streams for 1976­2000 on the basis of the OMNI cat alog of solar wind measurement data (see [Yermolaev et al., 2009] and the site ftp://ftp.iki.rssi.ru/pub/ omni/catalog/). The authors of a large number of papers treat dif ferent types of solar wind as isolated events and neglect
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Fig. 5. The SOHO/EIT image of the lower corona: the light regions correspond to active zones and the dark regions correspond to coronal holes [http://www.lmsal.com/YPOP/ProjectionRoom/latest/eit/full/eit284 128.gif].

their interaction. In most cases, this assumption is jus tified, because the sizes of the above listed disturbed types of phenomena in the Earth's orbit (1 AU) amount to no more than a few tenths of a percent of 1 AU and do not have time to significantly change or dissipate. However, interactions of CMEs in the vicin ity of the Sun [Gopalswamy et al., 2001, 2002] and of magnetic clouds in the vicinity of the Earth (see, for example [Burlaga et al., 2001; Berdichevsky et al., 2003; Gonzalez Esparza et al., 2004; Farrugia et al., 2006a] and references there in) were noted. In a num ber of papers, it is shown that some strong magnetic storms, for example, such as the events on March 31, 2001, with Dst = ­387 nT; on April 11­13, 2001, with Dst = ­271 nT [Wang et al., 2003]; on October 28­30, 2003, with Dst = ­363 nT [Veselovsky et al., 2004, Skoug et al., 2004]; on November 20, 2003, with Dst = ­472 nT [Yermolaev et al., 2005a, Gopalswamy et al., 2005]; and on November 8­10, 2004, with Dst = ­373 nT [Yermolaev et al., 2005b], were generated by interact ing magnetic clouds. Recently, in studying the inter planetary conditions for the 1995­2003 magnetic storms [Farrugia et al., 2006b], it was found that "a rather large number of our significant events (6 out of 16) consisted of interacting ICMEs/MCs that formed complex ejecta." Other authors [Xie et al., 2006] have studied 37 long lived geomagnetic storms observed in 1998­
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2002 with Dst < ­100 nT, which are associated with CMEs, and found that 24 of them were caused by a succession of CMEs and that the number of interact ing magnetic clouds was from 2 to 4. This result can be

6 The sun 2 4 3 5 4
Fig. 6. Schematic representation of the large scale types of solar wind: (1) heliospheric current sheet; (2, 3) slow and fast streams from coronal streamers, and coronal holes, respectively; (4) compressed plasma (at the boundary between fast and slow flows (CIR) and before the leading edge of a "piston" (sheath)); (5) "piston" (magnetic clouds (MCs) and "pistons" (ejecta)); and (6) rare plasma at the front of slow and fast flows of solar wind. Vol. 46 No. 7 2010

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Yu. I. YERMOLAEV, M. Yu. YERMOLAEV Chapman­Ferraro current electrojet

tail current

the Earth solar wind magnetopause

magnetopause current

ring current

tail current

Fig. 7. Structure of the Earth's magnetosphere. The solar wind compresses the geomagnetic field on the subsolar side. The magnetopause within which the Chapman­Fer raro current flows is formed where the pressures are equal. Both tail and ring currents are also shown.

Fig. 8. Scheme of electrojet formation during a magnetic substorm when electric current starts to flow through the Earth's ionosphere.

explained by a compression of magnetic clouds result ing in the formation of complex structures (complex ejecta) that simultaneously have the features of both magnetic cloud and a sheath (below, it will be shown that a sheath is more efficient in generating storms than the body of a magnetic cloud). 4. MAGNETOSPHERIC PHENOMENA The solar wind plasma and the plasma in the vicin ity of the Earth are practically ideal conductors of electric current. Therefore, in accordance with the laws of magnetic electrodynamics, the outer plasma of the solar wind and the IMF cannot closely approach the Earth because of its strong magnetic field. The interaction of the solar wind and the IMF with the Earth's plasma and magnetic dipole results in the for mation of a cavity (magnetosphere) at the boundary (magnetopause) of which the plasma and field (of both outer and inner origins) pressures are balanced. This magnetopause in the subsolar region is moved to a dis tance of about 6 â 104 km away from the Earth (Fig. 7). As a first approximation, the magnetosphere is impenetrable for the outer plasma of the solar wind, which can change only the form of the magnetopause in accordance with its pressure balance condition. However, in fact, the situation is more complicated. When the IMF has a component that is parallel to the Earth' magnetic dipole (the IMF southward compo nent), in the region of contact of the oppositely directed interplanetary and terrestrial magnetic fields, the condition of ideal plasma conductivity is violated and a magnetic field erosion occurs. The plasma of the solar wind and its transported energy enter the magnetosphere. This process is called the threshold (trigger) mechanism. According to the rate of energy

arrival, there are three possible scenarios of the mag netosphere's reactions. (1) When the rate of energy arrival is lower than or equal to the rate of stationary energy dissipation within the magnetosphere, its form does not change; no sig nificant changes are observed in the magnetosphere, i.e., the magnetosphere remains undisturbed. (2) When the rate of energy arrival exceeds that of stationary dissipation, a portion of energy leaves the magnetosphere through a "quasistationary channel," which results in its recovery. The role of such a channel is played by magnetic substorms (the processes of releasing magnetic energy accumulated in the mag netosphere through tail current connection (Fig. 7) along the magnetic lines through the ionosphere in the night region of the polar oval). The newly formed cur rent is called an "electrojet" (Fig. 8). The most impressive substorms--auroras--result from a bom bardment of atmospheric neutral atoms by accelerated (along the magnetic force lines) particles in the tail of the magnetosphere. The magnetosphere can, for a long time, release excess energy into the polar regions of both of the Earth's hemispheres in the form of sub storms with a periodicity of about 3 h. (3) When the rate of energy arrival significantly exceeds the rate of both stationary and quasistationary dissipations, global changes occur in both magneto spheric and ionospheric current systems, which are accompanied by strong disturbances of the Earth's magnetic field; this is called a magnetic storm. The basic contribution to the magnetic field's changes is made by the ring current located in the geomagnetic equator region (Fig. 7). Therefore, unlike magnetic substorms during which magnetic field disturbances are observed in the polar regions, during magnetic storms, the magnetic field varies also in low latitudes in the vicinity of the equator. During strong storms,
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180° 180° 220° 140° 260° 100° 300°340° 20° 60°

807

80°

Kp 9

60°

6
40° 20° 0° ­20° ­40°

3

0 ­300

­200

­100

­60°

0 Dst, nT

­80°

Fig. 10. Relation between the extreme values of the indices Kp and Dst for 611 magnetic storms with ­300 < Dst < ­60 nT for 1976­2000. The solid line denotes the approximation of the data presented [Yermolaev and Yermolaev, 2003b].

Fig. 9. Scheme of the location of two networks of ground based stations whose data are used in determining the indi ces Dst (circles) and Kp (triangles). The asterisks denote the location of the magnetic poles, and the solid line denotes the geomagnetic equator.

auroras can descend by 20°­30° to the equator of the polar regions and can be observed in the low latitudes, like, for example, on October 30, 2003. Thus, magnetospheric disturbances result from rapid changes in the current systems existing in the Earth's magnetosphere and ionosphere or from the formation of new current systems. It is important to note that the ring current variations during storms sig nificantly the exceed electrojets that occur during sub storms. However, because the ring current is located far from the Earth's surface (unlike the electrojet, which almost reaches the ionospheric and atmo spheric lower layers), during magnetic storms, varia tions in the Earth's magnetic field are of a global char acter (except for small regions in the vicinity of the magnetic poles) and amount to no more than 500 nT at the most. During substorms, the variations in the

magnetic field are of a local character and can amount to 1­3 â 103 nT (it should be remembered that the Earth's constant field amounts to about 30­50 â 103 nT; i.e., in any case, we are dealing with variations that do not exceed a few percent, which is significantly smaller than the fields of technogenic origin). The magnetospheric state is described by a number of different indices calculated on the basis of ground based measurements of the magnetic field [Mayaud, 1980]. Since the results obtained at different networks of magnetic stations are used to construct these indi ces, the latter include the responses of different mag netospheric and ionospheric current systems. Figure 9 shows the location of two networks of ground based stations whose data are used to calculate the indices Dst and Kp that are most often used to describe magnetic storms. On the one hand, it is possible to assume that, if the magnetic storm statistics are sufficient, there must be a correlation between the extreme values of different indices. Such an analysis was made for 1085 magnetic storms during the period 1957­1993 [Loewe and

Table 2. Classification of magnetic storms on the basis of the Dst index measured in 1957­1993 [Loewe and Prolss, 1997] Storm class Number of storms Weak Moderate Strong Severe Great 482 346 206 45 6 Percent 44 32 19 4 1 Dst, nT ­30...­50 ­50...­100 ­100...­200 ­200...­350 <­350 Dst ­36 ­68 ­131 ­254 427
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ap 27 48 111 236 300
2010

Kp 4 5 7 8 9
o o

AE, nT 542 728 849 1017 1335

­ + ­

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808

Yu. I. YERMOLAEV, M. Yu. YERMOLAEV () Flare Magnetic storm 150 Amount of flares 31.1% 100 (b) 48.2%

d 0

c

b

a 3 4

b 5

c

d 6 t, day

50

11.6% 9.0%

1 1.5 2

0 2000 1000 1500

500 400

300 V, km/s

0

a

b

c

d

Fig. 11. (a) Scheme of establishing correspondence between flares and storms for different time windows a, b, c, and d (the win dow length is given in days on the horizontal line, and the scale under this line shows the mean velocity of motion of disturbances from the Sun to the Earth in km/s. (b) The amount of western (gray columns) and eastern (white columns) strong solar flares that result in storms with a (a) high, (b) mean, or (c) low probability, as well as those that (d) do not result in storms. Their portion from the total amount of flares falling within the corresponding window is given in % [Yermolaev and Yermolaev, 2003a].

Prolss, 1997] (see Table 2). On the other hand, such results may create the illusion that the magnetospheric indices are interchangeable. However, even the first attempts to analyze real data show that the behavior of different indices is not identical during one and the same event. For example, on October 24, 2003, 15:00­23:00 UT [Veselovsky et al., 2004], at high val ues of Kp, the index Dst remained at undisturbed level. Figure 10 shows the relationship between the extreme values of the indices Kp and Dst for 611 mag netic storms (­300 < Dst < ­60 nT) for 1976­2000 [Yermolaev and Yermolaev, 2003b]. A wide scatter of data is explained by the fact that the indices Kp and Dst are measured in different geomagnetic latitudes and are sensitive to different current systems (magneto spheric phenomena): auroral electrojet (magnetic substorms) and ring current (magnetic storms). Thus, in order to study the relation between magnetic storms and different phenomena and to eliminate auroral phenomena from analysis, it is necessary to use the Dst index. In studying the influence of the auroral electro jet on different systems, it is better to use the special AE index. The index Kp is sensitive to both phenomena and does not allow the influence of each of the current systems to be individually studied. To correctly use the indices of geomagnetic activity in related disciplines (including heliobiophysics), it is necessary to have a general idea of the principles of constructing geomagnetic indices, their physical meaning, their interdependence, and the ranges of their values which correspond to different levels of geomagnetic activity. A familiarity with the data pre sented in this paper (Figs. 7, 8; Table 2) will provide additional insight into some of these problems. More detailed information can be found in [Mayaud, 1980; Loewe and Prolss, 1997].

5. CORRELATION BETWEEN DIFFERENT EVENTS OF "SPACE WEATHER" A lack of strong evidence of the cause­effect rela tion between the phenomena under study is customary in solar­terrestrial physics. The only experimental fact is that the one event was observed after the other during a time window previously specified. As a rule, any additional information on the phenomena is not direct in studying the relations between them. Let us consider the flare­magnetic storm relation as an example (Fig. 11). In this case, the time windows are determined by the mean speed of the movement of disturbances from the Sun to the Earth (the corre sponding scale is shown in Fig. 11a). The windows a, b, c, and d correspond to the intervals for which it is assumed that the flare­storm relation has a high, mean, low, and zero probability, respectively. Figure 11b shows the probability of storm generation after flares on the Sun's both western and eastern hemispheres. The total probability of storm generation by solar flares is assessed at about 40% [Yermolaev and Yermo laev, 2003a]. Only two levels of probability (related and unrelated events) are used in many studies, and, in this case, the fact that the relation between the events is of probabilistic character is neglected. The methods of identifying and classifying the solar (CMEs and solar flares), interplanetary (CIRs, sheaths, ejecta, and others), and geomagnetospheric (magnetic storms) phenomena have been described above. In addition to the ambiguity associated with different approaches to the classification of these phe nomena, there is an ambiguity associated with differ ent methods of comparing them in two spatial regions. If two phenomena with the sets X1 and X2 are chosen for analysis and the correspondences between these phenomena are found for the number of phenomena
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X12, the probability of the process X1 X2 is usually determined as the ratio X12/X1, which differs from the probability X2 X1 equal to X21/X2 = X12/X2. The values of X1 and X2 correspond to different phenom ena, are obtained according to different criteria, and can have different meanings (for example, it follows from Fig. 3 that the number of solar flares of class M and higher is approximately one order of magnitude larger than the number of moderate and strong mag netic storms). Hence, the geoeffectiveness determined in different studies depends on the direction of an analysis of the process. If we take into consideration that, in some studies, set X2 is not specified previously (before the start of analysis), i.e., the rules (or criteria) of choosing the phenomena for set X2 are not estab lished initially but are selected during the study, the ambiguity of calculating the geoeffectiveness can be increased additionally. Because in analyzing this chain of solar­terrestrial physics, a two stage process (the Sun­solar wind and the solar wind­magnetosphere) is studied, data on the interplanetary medium make it possible to increase the accuracy of the estimate of the whole chain. Let us assume that there are sets of data on the Sun (X1 and Y1), the interplanetary medium (Y2 and Z1), and the magnetosphere (X2 and Z2) for which the probabilities of the processes X1 X2 (X12/X1), Y1 Y2 (Y12/Y1), and Z1 Z2 (Z12/Z1) are estimated. In this case, it is natural to assume that the probability of the complete process must be close to the product of the probabilities of individual stages; i.e., X12/X1 = (Y12/Y1)(Z12/Z1). In particular, this implies that the geoeffectiveness (in the sense of probability) of the complete process cannot be higher than the geoeffec tiveness of each of the stages of X12/X1 Y12/Y1 and X12/X1 Z12/Z1. For example, the geoeffectiveness of solar events cannot exceed that of interplanetary events. Published data contain enough material to make such an analysis; its results will be demonstrated below. It is important to note that many authors often call values obtained with quite other methods the geoef fectiveness of phenomena. As was noted above, in the strict sense of the word, the geoeffectiveness (as a probability) of solar and interplanetary phenomena is a portion (percent) of the corresponding sets of solar and interplanetary phenomena resulting in a magnetic storm of certain intensity. The most common error is that some authors use the method of backward event tracing: first, they take the list of magnetic storms and then extrapolate them backward to the interplanetary medium or to the Sun in order to find corresponding phenomena. This method is important because it allows one to find phe nomena (candidates) that can be treated as the causes of the studied magnetic storms in the interplanetary medium and on the Sun. However, this method does not allow one to determine the geoeffectiveness (prob ability) of these phenomena. Phenomena of different
IZVESTIYA, ATMOSPHERIC AND OCEANIC PHYSICS

classes (if only they fall within the time window under analysis) are often treated as candidates, and this is one of the main reasons for disagreement among results obtained by many authors. As was noted in the Introduction, a study of the flare­magnetic storm relation and, generally, the rela tion between the events occurring on the Sun and in the interplanetary space is very important for heliobio physics, because geomagnetic disturbances on the Earth are predicted on its basis. This study gives infor mation necessary to develop preventive measures and, sometimes, to save patients with a cardiovascular pathology from fatal outcomes [Gurfinkel', 2004]. Improving the prediction of geomagnetic disturbances on the basis of solar observations, which will be dis cussed in the next section, is of fundamental impor tance for heliobiophysics. 6. RESULTS AND DISCUSSION With the preceding taken into account, data pub lished by different authors have been analyzed and some results on the geoeffectiveness of solar and inter planetary phenomena have been obtained in the con text of two different interpretations of geoeffective ness: as the probability of the generation of magnetic storms and as the efficiency of their generation. These results are discussed individually in the following sub sections. 6.1. Geoeffectiveness (Probability) of Different Phenomena The published results on the geoeffectiveness of CMEs, solar flares, and interplanetary phenomena were selected with consideration for event tracing (for ward or backward) and for the pairs of phenomena under analysis: CME storm; CME MC, ejecta; MC, ejecta storm; storm MC, ejecta; MC, ejecta CME; storm CME; flare storm; and storm flare. The results of this selec tion are given in Table 3 and schematically shown in Fig. 12. Table 3 differs from the previous publications [Yermolaev and Yermolaev, 2003b, 2006; Yermolaev et al., 2005a] in the inclusion of a number of recent papers and the additional process CIR storm [Alves et al., 2006]. The entry "CME storm" in Table 3 implies that the CME list with the numbers of analyzed cases (indicated in column 3) is used as an initial set of data. CMEs are compared with magnetic storms whose values are determined by the indices given in column 4 of Table 3. Thus, the data published were classified into six (I­VI) groups of phenomena comparison: three spatial regions and two directions of tracing (Table 3). Groups II, III, IV, and V combine magnetic clouds and ejecta (including sheaths and bodies) that are close to each other in origin and phys ical characteristics; however, in the column "The number of cases," symbols MC and E correspond to
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Yu. I. YERMOLAEV, M. Yu. YERMOLAEV

Table 3. Correlation between solar, interplanetary, and magnetospheric phenomena No. 1 1 2 3 4 5 6 7 8 9 10 11 12
a

% 2

Number of cases 3 I: CME storm 38 7 40 20 ? 132a 132a 125a 125a 70 b 49c 12c 218a 305b 229b

Index values, comments 4 Kp Dst Kp Kp Dst Kp Kp Dst Dst Dst Dst Dst Dst Dst Dst

Sources of data 5 Webb et al., 1996 Crooker, 2000; Webb et al., 2000; Li et al., 2001 Plunkett et al., 2001 Berdichevsky et al., 2002 Webb, 2002 Wang et al., 2002 Yermolaev and Yermolaev, 2003a Yermolaev and Yermolaev, 2003b Zhao, Webb, 2003 Moon et al., 2005 Yermolaev and Yermolaev, 2006 Kim et al., JGR, 2005 Gopalswamy et al., 2007

50 71 35 45 35­92 45 20 35 40 64 71 58 42 40 71

< > > < > > < < < < < < < <

­ 50 6 5 ­50 5 7 ­60 ­50 ­50 ­50 ­50 ­50 ­50 ­50

denotes halo CMEs directed toward the Earth, b denotes halo CMEs on the Sun's visible side, c denotes halo CMEs in the center of the Sun's visible side

II: CME 1 2 3 4

magnetic clouds, ejecta the earthward the earthward Cane et al., 1998 Webb et al., 2001 Berdichevsky et al., 2002 Schwenn et al., 2005

63 8 Halo CME toward 60­70 89 Halo CME toward 80 20 Halo CME 50­84 181 All CMEs All CMEs toward 53­90 154 59­93 91 Earthward full halo III: magnetic clouds, ejecta storm 44 67 63 57 19 63 82 73 50 43 76 56 327E 28 M C 30 M C 48 M C 1273E 1188E 34MC 135MC 214E 214E 104MC + E 104MC + E 104MC + E 149MC IV: storm CME 8 18 ? 27 23 10 54 K p > 5­ Dst < ­60 Dst < ­60 Dst < ­60

the earthward CME Gosling et al., 1991 Gopalswamy et al., 2000 Yermolaev and Yermolaev, 2002 Yermolaev et al., 2000 Gopalswamy et al., 2001 Yermolaev and Yermolaev, 2003b Richardson et al., 2001 Wu and Lepping, 2002a Wu and Lepping, 2002b Cane and Richardson, 2003 Zhang et al., 2004

1 2 3 4 5 6 7 8 9

Kp < 5­; Solar minimum Kp < 5­; Solar maximum Dst < ­50 Dst < ­50 Dst < ­50 Dst < ­60 D * < ­30 st D * < ­50 st D * < ­100 st Dst < ­50 Kp > 6 Kp > 6 ? Dst ­100 Dst ­100 ­100 < Dst < ­200 Dst < ­100

10 1 2 3 4 5 6

34 77 100 83 94 96 83 100 83

Echer and Gonzalez, 2004; Echer et al., 2005 Brueckner et al., 1998 St. Cyr et al., 2000; Li et al., 2001 Srivastava, 2002 Zhang et al., 2003 Watari et al., 2004 Srivastava and Venkatakrishnan, 2004
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SOLAR AND INTERPLANETARY SOURCES OF GEOMAGNETIC STORMS Table 3. (Contd.) 1 1 2 3 4 73 67 25 19 12 22 45 88 63 50 70 92 100 33 25 52 32 21 76 38 70 24 31 70 32 46 75 29 2 V: storm 37 12 1273E 833 352 62 26 1188E 670 341 99 78 618 414 204 90 100 21 21 30 150 58 10 187 63 8 271 VI: magnetic clouds, ejecta 1 2 3 4 5 6 1 2 3 4 1 1 2 3 1 3 magnetic clouds, ejecta Kp > 7 Dst < ­50 Dst (corr) Kp > 5­, Solar minimum 5 + > Kp > 5 ­ 6 + > K p > 6­ 7+ > Kp > 7­ K p > 8­ Kp > 5­, Solar maximum 5 + > Kp > 5 ­ 6 + > K p > 6­ 7+ > Kp > 7­ K p > 8­ Dst ­60 ­100 Dst ­60 Dst ­100 ­100 < Dst ­50 7 ­ > Kp > 5 ­200 Dst ­100 7 ­ > Kp > 8 Dst ­100 Dst ­50, 1978­1982 Dst ­100 Dst ­200 Dst ­50, 1995­2002 Dst ­100 Dst ­200 D * < ­30 st CME Lindsay et al., 1999 Cane et al., 2000 Gopalswamy et al., 2000 Burlaga et al., 2001 Cane and Richardson, 2003 Vilmer et al., 2003 Yermolaev and Yermolaev, 2002 Yermolaev and Yermolaev, 2003a Ivanov and Miletskii, 2003 Yermolaev and Yermolaev, 2006 Park et al., 2002 Krajcovic, Krivsky, 1982 Cliver and Crooker, 1993 Yermolaev and Yermolaev, 2003a Alves et al., 2006
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811

4

5 Gosling et al., 1991 Webb et al., 2000 Vennerstroem, 2001 Richardson et al., 2001 (GRL)

5

Yermolaev and Yermolaev, 2002

6

Huttunen et al., 2002

7 8

Watari et al., 2004 Li and Luhmann, 2004

9

Zhang et al., 2004

67 49E Any CME 65 86E Any CME 42 86E Earthward Halo CME 82 28MC Any CME 50­75 4MC Halo CME 40­60 5E Halo CME 56 193E Any CME 48 21MC Halo CME VII: flare storm 44 40 33 44 126 653 571 746 VIII: flare >M0 + SEP (Solar Energetic Particles) >M5 >3 (optics) >M5 SSC >M0 flare K p > 7­ Dst ­250 Dst ­100 storm 727 Dst ­50

35­45 4836 IX: storm 59 88 20 X: CIR 33 116 25 204

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812

Yu. I. YERMOLAEV, M. Yu. YERMOLAEV STORM Dst < ­50 Kp > 5

35­50% CME Magnetic cloud ejecta

60­80%

MC 60­80% E 40­80% STORM Dst < ­50 Kp > 5

P(CMEST) = 0.35­0.5 = P(CMEMC, E) · (MC, EST) = 0.3­0.6

80­100% CME Magnetic cloud ejecta

40­80%

30­70%

P(STCME) = 0.8­1.0 P(MC, EMCE) · P(STMC, E) = 0.1­0.6

Fig. 12. Schematic representation of the correlation between CME, MC/ejecta, and magnetic storms for the forward (top) and backward (bottom) tracings of the phe nomena. The numerical values of correlation are given above the corresponding arrows. The probabilities for both one and two step tracings are shown under each of the fragments [Yermolaev et al., 2005; Yermolaev and Yermo laev, 2006].

magnetic clouds and ejecta, respectively. Table 3 also includes data on the following groups of phenomena comparison: flare storm (group VII), flare storm sudden commencement (SSC) (group VIII), storm flare (group IX), and CIR storm (group X). An analysis of the publications on the CME storm process allows one to conclude that, for mag netic storms with Kp > 5 (Dst < ­50 nT), the geoeffec tiveness of halo CMEs directed toward the Earth amounts to 40­50% at sufficiently high statistics (from 38 to 305 CMEs) and the values obtained in [Webb, 2002; Zhao and Webb, 2003; and Gopalswamy et al., 2007] prove to be overestimated (see below). The results of an analysis of the backward tracing for group VI (storm CME) are in good agreement with each other (83­100%), but they reflect a low CME geoef fectiveness; they demonstrate that it is easy to find a probable candidate for the source of a storm among CMEs occurring on the Sun (this is not surprising, because the frequency of CME observations is a few times higher than that of magnetic storm recordings). An analysis of the sequence of the two step forward tracing for groups II (CME magnetic clouds, ejecta) and III (magnetic clouds, ejecta storm)

makes it possible to estimate the probability of the complete CME storm process as a product prob ability, and, for magnetic clouds, we obtain (0.60­ 0.70)(0.57­0.82) = 0.34­0.57, which is close to the above mentioned result (40­50%) for the one step CME storm process (group I) and is lower than the estimates obtained in [Zhao and Webb, 2003; Gopalswamy et al., 2007]. For ejecta, this approach yields a smaller value. An analysis of the two step backward tracing for groups V (storm magnetic clouds, ejecta) and VI (magnetic clouds, ejecta CME) does not allow a high value (83­100%) to be obtained for the complete process (storm CME), because, in this case, the product probability yields (IV) (0.25­0.73)(0.42­0.82) = 0.11­0.60. Thus, the results of comparison between the one and two step processes for the forward tracing (CME storm) are in good agreement, while, for the backward trac ing, the two step process is a few times smaller than the one step process. This implies that the values of the processes (storm magnetic clouds, ejecta), (magnetic clouds, ejecta CME), and (storm CME) are not the geoeffectivenesses (probability) of the processes describing the "cause effect" sequence. Although the "effectiveness" (probability) of storm generation in [Webb et al., 2000; Webb, 2002; Zhao and Webb, 2003; Gopalswamy et al., 2007] refers to the forward tracing of I (CME storm) and is lower than that for the backward tracing of IV (storm CME), the results obtained in the indicated papers (1) exceed estimates obtained in other studies of this process; (2) exceed the values obtained for the second step, i.e., for process III (except for [Wu and Lepping, 2002a,b; Echer and Gonzalez, 2004; Echer et al., 2005]); (3) are close to the value obtained for the first step, i.e., process II; and (4) exceed the estimates obtained for two step process II: III = (0.6­0.8)(0.2­0.8) = 0.1­ 0.6. Thus, the geoeffectiveness estimates obtained in [Webb et al., 2000; Webb, 2002; Zhao and Webb, 2003; Gopalswamy et al., 2007] are apparently overesti mated. As was noted above (Section 2), there is a prob ability that some CMEs directed toward the Earth could inaccurately be referred to the CMEs on the Sun's back side [Zhang et al., 2003]; it is possible that this methodical error was made in analyzing data from a number of individual papers [Yermolaev, 2008]. In our recent studies, we [Yermolaev and Yermo laev, 2002, 2003a] performed the forward tracing of the events (flare storm) and assessed the geoeffective ness (probability) of 653 strong solar flares of the X ray class higher than M5 and that of 126 weaker solar flares of the X ray class higher than M0, but which were accompanied by an increasing flux of energetic parti cles in the vicinity of the Earth. For magnetic storms with Dst < ­60 nT the geoeffectiveness (probability) , amounted to 40% in the former case and 44% in the latter case. If backward tracing (storm flare) is performed and a list of strong magnetic storms with Dst < ­100 nT is
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taken, only 20% of the indicated flares can be the sources of such storms. Similar investigations were carried out earlier. For example, in [Krajcovic and Krivsky, 1982], the backward tracing (storm flare) was analyzed for strong solar flares of the optical class and it was shown that, in 1954­1976, 59% of possible sources were found for 116 storms with Kp > 7. In [Cliver and Crooker, 1993], the backward tracing (storm flare) showed that, for the 25 strongest magnetic storms (Dst < ­250 nT) during 1957­1990, flares that could be treated as candidates for storm sources were found in at least 22 cases (88%).The high values of geoeffectiveness (probability) obtained in [Krajcovic and Krivsky, 1982; Cliver and Crooker, 1993], in addition to the backward tracing of events, can be associated with the fact that, in these papers, even weak flares were treated as possible candidates for storm sources (see the ratio between the frequencies of flares and storms in Fig. 3), while only strong flares were used in our analysis. A comparison of the (flare SSC) events (i.e., not with magnetic storms but with their precursors), which are often observed a few dozen minutes before the start of the main phase of a magnetic storm, was made in [Park et al., 2002] for 4836 flares of the class higher than M1 for 1975­1999. The results of this comparison showed that the estimate of geoeffective ness (probability) amounts to 35­45% for a time delay (waiting "window") of 2­3 days and 50­55% for a longer window. This result is close to the above dis cussed geoeffectiveness of flare storm, although it was obtained for SSCs but not for storms. To assess the probabilities from the practical point of view, let us compare the probabilities of magnetic storm generations by solar and random events. To this end, if we take the value of the waiting window and the mean period between storms during disturbed years, we will obtain 35% [Yermolaev and Yermolaev, 2003b]. This implies that, currently, the probability of predict ing magnetic storms on the basis of solar data only slightly exceeds the probability of prediction based on randomness. Thus, the reliability of predicting magnetic storms on the basis of interplanetary phenomena and direct measurements of solar wind parameters (see, for example, http://www.iki.rssi.ru/sw.htm) is sufficiently high (60­80% and 90­95%, respectively). In addi tion, the parameters of the solar wind and the IMF make it possible to calculate the values of magneto spheric disturbances. The reliability of prediction on the basis of solar phenomena is low; a predictable magnetic storm can be estimated according to the class of solar flare only after the record of its associated CME moving toward the Earth (i.e., halo CME). For example, the close algorithms of predicting magnetic storms on the basis of observations of solar flares and CMEs (or CME related preceding shock wave and radio bursts of types IV and II, respectively) are pro posed in [Song et al., 2006; and Valach et al., 2007].
IZVESTIYA, ATMOSPHERIC AND OCEANIC PHYSICS

The prediction made without consideration for CME data is basically of a random character and of no practical significance. The quality of such a prediction should be improved in two directions: (1) to take into consideration a quantitative description of interrela tions between the parameters of flare and CME and (2) to increase the probability of predicting interac tions between CME and the Earth's magnetosphere. 6.2. Efficiency of Storm Generation by Different Phenomena As was noted above, magnetic storms are generated mainly by the following types of solar wind: ICME, including sheath and the ICME body (MC/ejecta), and CIR [Vieira et al., 2004; Huttunen and Koskinen, 2004; Yermolaev et al., 2005c; Yermolaev and Yermo laev, 2006]. In [Yermolaev and Yermolaev, 2002], it is shown that the time variations in the portion of storms induced by CIR and ICME have two maxima (min ima), each over the solar cycle and change in antiphase. On the other hand, there are experimental data on differences in storms generated by sheath, MC, and CIR [Borovsky and Denton, 2006; Denton et al., 2006; Pulkkinen et al., 2007b]. The ratios between extreme values (peak to peak analysis) Bz­ Dst and Ey­Dst (Ey = VxBz is the electric field of the solar wind during the IMF southward Bz component) are given in some papers individually for CIR and MC induced storms. When considering these data simultaneously (Figs. 13, 14), one can see that, against the background of an evident scatter in experimental points, there are no significant differences in depen dencies for different sources of storms. These results were obtained without selecting sheaths and MCs, and since the conditions of sheaths and CIRs are close to each other (both types are formed due to the compres sion and deformation of slow and fast flows), this approach could mask differences in the dependences under consideration. To verify if there is such a possi bility, we calculated the dependences Bz­Dst and Ey­Dst individually for CIR, sheath, and MC; however, against the background of wide data scattering, we also found no significant differences [Yermolaev et al., 2007b]. Therefore, differences in storms induced by CIRs, sheaths, and MCs can be associated not with the peak values of Bz and Ey, but with other solar wind plasma and field parameters and their dynamics. For 623 magnetic storms with Dst < ­60 nT, which were recorded in 1976­2000, we analyzed the effects of the solar wind and IMF parameters and their vari ations individually for CIRs, sheaths, and MCs on the basis of the OMNI database (with calculated addi tional physical parameters) [Yermolaev et al., 2007a, b]. For this analysis, we used the method of the superpo sition of epochs (the start time of storm was taken as the start of the epoch). The differences in the time profiles of the solar wind and IMF parameters for CIR (121 storms), sheath (22 storms), and MC
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814 Dst, nT 100 0

Yu. I. YERMOLAEV, M. Yu. YERMOLAEV Dst, nT 100 0 1 ­100 ­200 ­300 ­400 4 8 5 6 ­400 ­500 0 4 8 12 16 22 3 2 ­200 ­300 3 4 1 2 5 6 7 26 30 Ey, mV/m 7 ­100

­500 ­70 ­60 ­50 ­40 ­30 ­20 ­10 0 Bz, nT
Fig. 13. Correlation dependences between the index Dst and the Bz component for different types of solar wind according to data obtained by different authors for CIR (1) [Alves et al., 2006], (2) [Richardson and Cane, 2005], and (3) [Yermolaev et al., 2007b], as well as for sheath + MC (4) [Naitamor, 2005], (5) [Wu and Lepping, 2002b, 2005], (6) [Richardson and Cane, 2005], (7) [Yermolaev et al., 2007b], and (8) [Yurchyshyn et al., 2004].

Fig. 14. Correlation dependences between the index Dst and the Ey component for different types of the solar wind according to data obtained by different authors for CIR (1) [Alves et al., 2006] and (2) [Yermolaev et al., 2007b] and for sheath + MC (3) [Srivastava and Venkatakrishnan, 2004], (4) [Kane, 2005], (5) [Wu and Lepping, 2005], (6) [Yermolaev et al., 2007b], and (7) [Wu and Lepping, 2002b].

(113 storms) are shown in Fig. 15, where 367 storms for which full datasets were absent in the OMNI data base (about half the total observation period) are denoted as "unknown," which did not allow the types of the solar wind to be unambiguously identified. In the left column, the following parameters are given: density N, velocity V, dynamic pressure Pdyn, proton temperature T, ratio between measured and calculated (on the basis of velocity) temperatures T/Texp, and the index Dst. In the left column, the following parameters are presented: the ratio between thermal and magnetic pressures ; B, Bx, By, and Bz--the magnitude and geocentric solar magnetospheric (GSM) components of the IMF; and the index Kp. The curves obtained from averaging over different types of solar wind differ in color. The variability of all the parameters for differ ent types of the solar wind is sufficiently high, and, in some case, the differences between the points of the averaged curves are less noticeable than the variances of the corresponding parameters, which suggests that this is a tendency rather than a proven fact. Neverthe less, a noticeable divergence of the curves during mea surements for many hours suggests a number of defi nite conclusions [Yermolaev et al., 2007a, b]. (1) Although the behavior of the solar wind param eters during magnetic storms is markedly different for different types of solar wind, the IMF Bz component for all wind types turns southward 1­2 h before the start of magnetic storm and, 2­3 h after its start, reaches its minimum; simultaneously, both density and dynamic pressure increase. (2) Although the IMF Bz component's lower values are observed in MCs, the smallest mean of the index Dst is reached in sheaths. Thus, the generation of a

storm is more efficient during sheaths than during the arrival of the MC's body probably due to the higher values and variations of the pressure and magnetic field. (3) The fact that nkT, T/Texp, and parameter in CIR and sheath have higher values than in MC is in agreement with the physical nature of these types of solar wind and supports the correctness of their selec tion. 7. CONCLUSIONS At first, let us briefly review the basic principles of solar­terrestrial physics related to the sources and causes of magnetic storms on Earth. (1) The source of magnetic storms on Earth is the large (>5 nT) and long lasting (more than 2 h) south ward (Bz < 0) component of the IMF, which makes the magnetosphere "open" for an enter of solar wind energy. (2) In the stationary solar wind, the Bz component is small or quite absent; therefore, all magnetic storms are associated with the disturbed types of the solar wind. (3) There are two chains of solar­terrestrial relations which result in magnetic storms: (1) CME M C + it s preceding compression region (sheath) with Bz < 0 magnetic storm, and (2) coronal hole high solar wind forms a compression region with the Bz < 0 magnetic storm. (4) Only the solar flares that are associated with CMEs can be treated as candidates for the solar sources of magnetic storms. Most flares do not have any cause­effect relation with magnetic storms.
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SOLAR AND INTERPLANETARY SOURCES OF GEOMAGNETIC STORMS 100 1.0 N 10 600 V, km/s B, nT 500 0.1 20 15 10 5 10 Bx, nT 0 ­10 10 By, nT 0 ­10 10 Bz, nT 0 ­10 0 Dst, nT Kp · 10 ­40 ­80 ­120 ­12 ­6 0 6 12 18 24 80 60 40 20 ­12 ­6 0 6 12 18 24 h

815

400 Pdyn, nPa T, K T/T
exp

10

1

1E + 5

1E + 4

1

Fig. 15. Behavior of the parameters of the plasma and magnetic field of the solar wind for magnetic storms generated by different types of solar wind during 1976­2000: CIR (green curves), sheath (red curves), MC (blue curves), and "unknown" (black curves) [http://arxiv.org/abs/physics/0603251; Yermolaev et al., 2007]. The epoch superposed analysis method with the "zero" time cor responding to the first one hour point of a rapid decrease in the Dst index was used for an analysis of the OMNI database. The time from the start of epoch is plotted on the horizontal axes, h. (See the text to identify the parameters in both left and right col umns.)

(5) With an increase in the negative Bz component or energy arriving in the magnetosphere, magnetic substorms (auroral electrojets) occur first (at lower values) and then (at higher values) magnetic storms (ring current disturbances) occur together with sub storms. Auroral electrojets exert influence locally (in the auroral night region), and storms exert influence globally (within a broad band outside the auroral regions). (6) The most faithful indicator of magnetic storms on Earth is the index Dst; the indicator of substorms is the AE index. The index Kp is sensitive to both storms
IZVESTIYA, ATMOSPHERIC AND OCEANIC PHYSICS

and substorms and does not allow their influences to be separated. (7) The term "geoeffectiveness" has two different meanings: (1) probability, i.e., a portion (percent) of one or another of the phenomena that have the cause­ effect relations with magnetic storms (in this case, it is necessary to use the methods of the forward tracing from the phenomenon to the storm and not the reverse); and (2) the efficiency of storm generation by different phenomena that have the cause­effect rela tions with storms, i.e., a comparison between the "output" and "input" of the process.
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Yu. I. YERMOLAEV, M. Yu. YERMOLAEV M. V. Alves, E. Echer, and W. D. Gonzalez, "Geoeffective ness of Corotating Interaction Regions as Measured by Dst Index," J. Geophys. Res. 111 A07S05, doi: 10.1029/2005JA011379 (2006). D. B. Berdichevsky, C. J. Farrugia, B. J. Thompson, et al., "Halo­Coronal Mass Ejections near the 23rd Solar Minimum: Lift off, Inner Heliosphere, and in situ (1 Au) Signatures," Ann. Geophys. 20, 891 (2002). D. B. Berdichevsky, C. J. Farrugia, R. P. Lepping, et al. "Solar­Heliospheric­Magnetospheric Observations on 23 March­26 April 2001: Similarities to Observations in April 1979," in Solar Wind 10, AIP Conference Pro ceedings, Ed. by M. Velli (Woodbury, New York, 2003). J. E. Borovsky and M. H. Denton, "Differences between CME Driven Storms and CIR Driven Storms," J. Geophys. Res. 111 A07S08, doi: 10.1029/2005JA011447 (2006). V. Bothmer, "The Solar and Interplanetary Causes of Space Storms in Solar Cycle 23," IEEE Transact. Plasma Sci. 32 (4), 1411 (2004). G. E. Brueckner, J. P. Delaboudiniere, R. A. Howard, et al., "Geomagnetic Storms Caused by Coronal Mass Ejections (CMEs): March 1996 through June 1997," Geophys. Res. Lett. 25, 3019 (1998). L. F. Burlaga, R. M. Skoug, and C. W. Smith, "Fast Ejecta During the Ascending Phase of Solar Cycle 23: ACE observations, 1998­1999," J. Geophys. Res. 106, 20957 (2001). H. V. Cane and I. G. Richardson, "Interplanetary Coronal Mass Ejections in the Near Earth Solar Wind during 1996­2002," J. Geophys. Res. 108 (A4), 1156, doi: 10.1029/2002JA009817 (2003). H. V. Cane, I. G. Richardson, and O. C. St. Cyr, "The Interplanetary Events of January­May, 1997, as Inferred from Energetic Particle Data, and Their Rela tionship with Solar Events," Geophys. Res. Lett. 25 (14), 2517 (1998). H. V. Cane, I. G. Richardson, and O. C. St. Cyr, "Coronal Mass Ejections, Interplanetary Ejecta and Geomag netic Storms," Geophys. Res. Lett. 27 (21), 3591 (2000). E. W. Cliver and N. U. Crooker, "A Seasonal Dependence for the Geoeffectiveness of Eruptive Solar Events," Solar Phys. 145, 347 (1993). E. W. Cliver and H. S. Hudson, "CMEs: How do the Puzzle Pieces Fit Together?," J. Atmos. Solar­Terrestr. Phys. 64, 231 (2002). N. U. Crooker, "Solar and Heliospheric Geoeffective Dis turbances," J. Atmos. Solar­Terrestr. Phys. 62 1071 (2000). M. H. Denton, J. E. Borovsky, R. M. Skoug, et al., "Geo magnetic Storms Driven by ICME and CIR Domi nated Solar Wind," J. Geophys. Res. 111, A07S07, doi: 10.1029/2005JA011436 (2006). E. Echer and W. D. Gonzalez, "Geoeffectiveness of Interplan etary Shocks, Magnetic Clouds, Sector Boundary Cross ings and Their Combined Occurrence," Geophys. Res. Lett. 31, L09808, doi: 10.1029/2003GL019199 (2004). E. Echer, M. V. Alves, and W. D. Gonzalez, "A Statistical Study of Magnetic Cloud Parameters and Geoeffec tiveness," J. Atmos. Solar Terrestr. Phys. 67, 839­852 (2005).
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An analysis of publications on the geoeffectiveness of solar and interplanetary phenomena and the results obtained allowed us to conclude the following: (i) The geoeffectiveness estimate depends on the methods of identifying and classifying phenomena and on the methods and directions of searching for corre lations between phenomena (backward tracing does not yield geoeffectiveness estimates). (ii) The geoeffectiveness (the probability of mag netic storm generation) for CME and flares amounts to 40­60%, which only slightly exceeds the probabil ity of random processes. (iii) The prediction of magnetic storms on the basis of solar observations may contain a substantial per centage of "false alarms." (iv) The geoeffectiveness of ICME (sheath + MC) amounts to 60­80%. (v) The geoeffectiveness of CIR amounts to 20­ 35%. (vi) No significant differences were found in the peak to peak dependences of Dst­Bz and Dst­Ey for magnetic storms generated by MC, Sheath, and CIR, although there are differences in their development. (vii) The minimum Bz component of the IMF is observed in MC, and the minimum Dst index is observed in sheath; i.e., the efficiency of the physical process of storm generation during sheath is higher than during MC. This is possibly due to the higher level of field and pressure variations for sheath. ACKNOWLEDGMENTS We are grateful to researchers who work with SOHO, GOES, and other satellites, as well as at ground based stations, and those who developed the OMNI database for the opportunity to use these data in our studies. This study was supported by the Russian Founda tion for Basic Research (project no. 07 02 00042) and by the Department of Physical Sciences, Russian Academy of Sciences, Program 15 "Plasma Processes in the Solar System." REFERENCES
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