Linking two consecutive non-merging magnetic clouds with...

JOURNAL OF GEOPHYSICAL RESEARCH, VOL. ???, XXXX, DOI:10.1029/,

Linking two consecutive non-merging magnetic clouds with theirsolar sourcesS. Dasso1,2, C.H. Mandrini1, B. Schmieder3, H. Cremades4, C. Cid5, Y. Cerrato5,

E. Saiz5, P. Demoulin3, A.N. Zhukov6,7, L. Rodriguez6, A. Aran8, M. Menvielle9,and S. Poedts10

Abstract. On May 15, 2005, a huge interplanetary coronal mass ejection (ICME) wasobserved near Earth. It triggered one of the most intense geomagnetic storms of solarcycle 23 (Dstpeak = -263 nT). This structure has been associated with the two-ribbonflare, filament eruption and CME originating in active region (AR) 10759 (NOAA num-ber). We analyze here the sequence of events, from solar wind measurements (at 1 AU)and back to the Sun, to understand the origin and evolution of this geoeffective ICME.From a detailed observational study of in situ magnetic field observations and plasmaparameters in the interplanetary (IP) medium, and the use of appropriate models, wepropose an alternative interpretation of the IP observations, different to those discussedin previous studies. In our view, the IP structure is formed by two extremely close con-secutive magnetic clouds (MCs) that preserve their identity during their propagation throughthe interplanetary medium. Consequently, we identify two solar events in Hα and EUVwhich occurred in the source region of the MCs. The timing between solar and IP events,as well as the orientation of the MC axes and their associated solar arcades are in goodagreement. Additionally, interplanetary radio type II observations allow the tracking ofthe multiple structure through inner heliosphere and to pin down the interaction regionto be located midway between the Sun and the Earth. The chain of observations fromthe photosphere to interplanetary space is in agreement with this scenario. Our anal-ysis allows the detection of the solar sources of the transients and explains the extremelyfast changes of the solar wind due to the transport of two attached —though non-merging—MCs which affect the magnetosphere.

1. Introduction

Coronal mass ejections (CMEs) remove plasma and mag-netic field from the Sun expelling them into interplanetaryspace. When they are detected in the interplanetary (IP)medium, they are called interplanetary coronal mass ejec-tions (ICMEs). The transit time of an ICME from the Sunto 1 AU is typically in the interval 1 to 5 days [Gopalswamyet al., 2000, 2001a; Rust et al., 2005]. A subset of ICMEs,called magnetic clouds (MCs), are characterized by an en-hanced magnetic field strength, a smooth and large rotationof the magnetic field vector, and low proton temperature

1Instituto de Astronomıa y Fısica del Espacio (IAFE),CONICET-UBA, Buenos Aires, Argentina

2Departamento de Fısica, FCEN-UBA, Buenos Aires,Argentina

3Observatoire de Paris, LESIA, UMR 8109 (CNRS),Meudon, France

4Universidad Tecnologica Nacional, Facultad RegionalMendoza, Argentina

5Universidad de Alcala, Alcala de Henares, Spain6Solar-Terrestrial Center of Excellence - SIDC, Royal

Observatory of Belgium, Brussels, Belgium7Skobeltsyn Institute of Nuclear Physics, Moscow State

University, Moscow, Russia8Universitat de Barcelona, Barcelona, Spain9Centre detude des Environnements Terrestre et

Planetaires, Velizy, France10K.U. Leuven, Leuven, Belgium

Copyright 2008 by the American Geophysical Union.0148-0227/08/$9.00

[Burlaga et al., 1981]. Fast and large magnetic clouds aremostly observed at 1 AU in the declining phase of a solarcycle, as it happened in the last one during 2003-2005 [e.g.,Culhane and Siscoe, 2007].

Interplanetary type II bursts permit the tracking of fastICMEs through the analysis of radio emission frequency drift[Reiner et al., 1998, 2007; Hoang et al., 2007]. Assuming thetype II emission to be produced at the fundamental or sec-ond harmonic of the local plasma frequency, and with helpof a heliospheric density model, the radio frequency can beconverted to radial distance. This method has still limita-tions and allows for possible different interpretations due tothe patchiness and the frequency range of the observed radioemissions. The evolution of an ICME can also be followedalong its journey using a combination of solar and IP obser-vations together with the results of numerical simulations[e.g., Wu et al., 1999].

Coronal mass ejections are frequently associated with fil-ament eruptions. The directions of the MC axes are foundto be roughly aligned with the disappearing filaments [Both-mer and Schwenn, 1994, 1998], preserving their helicity sign.This result has been found for some individual cases byMarubashi [1997]; Yurchyshyn et al. [2001]; Ruzmaikin et al.[2003]; Yurchyshyn et al. [2005]; Rodriguez L. et al. [2008].Some quantitative (quantifying magnetic fluxes and helici-ties) studies of MCs and their solar sources have been alsodone [Mandrini et al., 2005; Luoni et al., 2005; Longcopeet al., 2007]. But again, it is not so easy to quantify thisassociation and further developments are needed to reallyunderstand the fundamental transport mechanisms and in-teraction with the ambient solar wind (see, e.g., the reviewby Demoulin [2008]).

Furthermore, when multiple CMEs are expelled from theSun, they can be merged leading to the so-called “CME

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X - 2 DASSO ET AL.: TRACKING TWO CONSECUTIVE MAGNETIC CLOUDS

cannibalism” [Gopalswamy et al., 2001b], mainly as a conse-quence of magnetic reconnection. The interaction of twoCMEs in favorable conditions for reconnection has beenstudied by Wang et al. [2005]. However, from a theoreticalpoint of view, two original structures can be preserved withan appropriate orientation of the ejected flux rope (yieldingalmost parallel interacting magnetic fields). For example, inthe MHD simulations of Xiong et al. [2007] the two inter-acting flux ropes preserved their identity while evolving inthe IP medium until they reach the Earth environment oreven beyond.

In the magnetohydrodynamic (MHD) framework, a mag-netic cloud configuration in equilibrium can be obtainedfrom the balance between the magnetic Lorentz force andthe plasma pressure gradient. Several magnetostatic modelshave been used to describe the configuration of MCs. Fre-quently, the magnetic field of MCs has been modeled by theso-called Lundquist’s model [Lundquist , 1950], which consid-ers a static and axially symmetric (cylindrical) linear force-free magnetic configuration neglecting the plasma pressure[e.g., Goldstein, 1983; Burlaga, 1988; Lepping et al., 1990;Lynch et al., 2003; Dasso et al., 2005; Leitner et al., 2007;Xiong et al., 2007]. The azimuthal and axial field compo-nents of the flux rope in this classical configuration are de-fined by Bφ = B0J1(αr) and Bz = B0J0(αr), where r isthe distance to the MC axis, Jn is the Bessel function ofthe first kind of order n, B0 is the strength of the fieldat the cloud axis, and α is a constant associated with thetwist of the magnetic field lines. Some other refined modelshave also been proposed to describe the magnetic structureof clouds [e.g., Hu and Sonnerup, 2001; Vandas and Ro-mashets, 2002; Cid et al., 2002]; in particular, some of themconsider an elliptical shape [Hidalgo et al., 2002a; Hidalgo,2003] that allows the description of possible distortions ofthe structures.

On May 15-17, 2005, the strongly southward interplan-etary field (above 40 nT) and the high solar wind velocity(close to 1000 km s−1), observed by the Advanced Com-position Explorer (ACE), are the cause of a super geomag-netic storm with a depression of the Dst index reaching-263 nT. The IP structure has the characteristics of a MC.Yurchyshyn et al. [2006] associated this structure with thetwo-ribbon M8.0 X-ray class flare on May 13 at 16:32 UT inAR 10759, accompanied by a filament eruption and CME.These authors identified a single MC, starting at ∼ 00:00 UTon May 15 and ending at ∼ 09:00 UT on May 16. Anevent with different boundaries, much smaller than the onedescribed by Yurchyshyn et al., was identified by R. Lep-ping (start on May 15 at 05:42 UT and end on May 15 at22:18 UT). The event was qualified with intermediate fittingquality (quality 2) and had a large impact parameter (cata-log at http://lepmfi.gsfc.nasa.gov/mfi/mag cloud S1.html).

The solar activity evolution along May 13 has been welldescribed by Yurchyshyn et al. [2006] and, more recently,by Liu et al. [2007]. Both papers refer mainly to the M8.0event and conclude that the eruption of a large sigmoidalstructure launches the CME observed by the Large Angleand Spectroscopic Coronagraph [LASCO, Brueckner G. E.et al., 1995] at 17:22 UT.

In this paper, we present a detailed analysis of the solarwind conditions at the Lagrangian L1 point using in situobservations (magnetic field and bulk plasma properties)and propose an alternative interpretation of the IP obser-vations. In our view, the observed ICME is in fact formedby two MCs. Since a single solar event cannot explain thearrival of two MCs at 1 AU, we revisit the evolution andactivity of AR 10759 and other regions present in the solardisk. We start searching in the Sun for two possible sourcesof the clouds, even earlier than May 13. We identify a pre-vious candidate event in Hα and EUV data same day at12:54 UT in AR 10759, which was classified as a C1.5 flare

in GOES. These two events, the one at 12:54 UT and theone at 16:32 UT, give rise to two two-ribbon flares alongdifferent portions of the AR magnetic inversion line. Theorientations of the magnetic fields associated to both solarevents are in good agreement with the fields observed intheir associated MCs.

In Section 2 we analyze the solar wind conditions nearEarth; in particular, we derive the orientation of the axisfor the two MCs by fitting a model to the observations. InSection 3, we revisit the solar events using ground basedand satellite data. In Section 4 we present radio type IIremote observations, which let us track the multiple struc-ture through the inner heliosphere. Finally, in Section 5,we describe the most plausible physical scenario behind thechain of events from May 13 to 17, 2005, and we give ourconclusions.

2. A Large Interplanetary Coronal MassEjection: May 15-17.2.1. In situ plasma and magnetic field observations

The solar wind data were obtained by ACE, located atthe Lagrangian point L1. The magnetic field observationscome from the Magnetic Fields Experiment (MAG) [Smithet al., 1998] and the plasma data from the Solar Wind Elec-tron Proton Alpha Monitor (SWEPAM) [McComas et al.,1998].

Figure 1 shows the magnetic field (in the Geocentric SolarEcliptic, GSE, system) and plasma conditions in the solarwind for the long-lasting studied event (00:00 UT on May15 to 16:00 UT on May 17, 2005), and the consequent mag-netospheric activity (Dst index).

The data in Figure 1 show several features consistent withthe presence of ICMEs/MCs: low proton plasma beta, pro-ton temperature lower than expected (Tex) for a typical so-lar wind at a given observed bulk velocity [Lopez , 1987], analmost linear velocity profile, and a smoothly varying mag-netic field orientation of high intensity. In particular, fromvertical dashed line ’1’ to ’10’ (for timings of tick numberssee next paragraphs and Table 1) the magnetic field inten-sity is strongly enhanced, although the scale needed to showthe whole event in one plot does not allow to observe thisenhancement at the end of the time interval. The tempo-ral length of the event is quite large, so that it corresponds

Table 1. Timings (dd,hh:mm UT) and substructures insidethe ICME of May 15-17, 2005,their boundaries are identifiedwith numbers in Figure 1.

Tick number Timing Substructure1 15,02:11

sheath2 15,05:42

MC13 15,10:20

back14 15,14:10

MC25 15,21:40

MC26 16,00:15

MC27 16,04:10

back28 16,13:00

back29 17,09:37

back210 17,10:30

DASSO ET AL.: TRACKING TWO CONSECUTIVE MAGNETIC CLOUDS X - 3

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Figure 1. In situ (at L1) plasma and magnetic field observations (ACE) of the ICME on May 2005.

From top to bottom: absolute value of the magnetic field (B = | ~B|), magnetic field vector orientation(GSE): latitude (θB) and longitude (φB), bulk velocity (V ), expected (continuous line) and observed(dots) proton temperature (Tp), proton density (np), proton plasma beta (βp), and Dst index. Verticallines mark different interfaces (see Section 2.1 for a description and Table 1 for timings). Horizontaldotted lines in θB , φB , βp, and Dst mark values at 0, 180, 0.1, and -100 as a reference, respectively.

to a very extended region, ∼ 0.93 AU, one of the largestICMEs ever observed (see, e.g., Table 1 of Liu et al. [2005]and Figure 1 of Liu et al. [2006a]). At ’10’, the magneticfield recovers its background value of ∼ 5 nT.

From a comparative study between magnetic clouds andcomplex ejecta, Burlaga et al. [2001] found that while mostof the clouds were associated with single solar sources, nearlyall the complex ejecta could have had multiple sources. Evi-dence was also found indicating that long duration complexmerged interaction regions (with radial extents of ∼ 0.7AUs) can be produced by the interaction of two or moreCMEs/MCs/shocks [Burlaga et al., 2001, 2003].

In front of the ICME, from 00:00 UT to 02:11 UT onMay 15 (labeled as ’1’), the solar wind presents typical con-ditions with a value of B ∼ 6 nT, an increasing velocity pro-file (starting at ∼ 400 km s−1 and reaching ∼ 500 km s−1),observed proton temperature similar to that expected one(Tex) for a typical solar at same velocity, an increasing pro-ton density profile from 3 to 10 protons per cm3, high valuesof proton plasma βp (reaching ∼ 10-100), and values of theDst index (Dst ∼ 0) corresponding to relatively quiet ringcurrent conditions.

Observations at ’1’ suggest the existence of a strong lead-ing edge shock (not fully confirmed from observations of V

because of a gap in plasma data from 02:11 UT to 03:50 UTon May 15). Just behind the shock, from ’1’ to ’2’, a typicalICME sheath is present (e.g., high level of fluctuations in ~B,enhanced B strength, high mass density, high βp).

A structure with a very high B (50-60 nT), one of thehighest values ever observed in the solar wind at 1 AU, isfound between ’2’ and ’3’. In this range of time, a large scalecoherent rotation of the magnetic field vector is present, with~B going from south to north (see θB panel). This is a lefthanded SEN flux rope in the classification of Bothmer andSchwenn [1998]. Note the presence of a sudden change ofthe sense of rotation together with a magnetic discontinu-ity (MD) at ’3’ (clearly observed in θB). A current sheet(i.e., a discontinuity in the observed time series of the mag-netic field vector) is expected to be present at the interfacethat separates two magnetic regions with different magneticconnectivity, i.e., with different magnetic stress. Thus, weinterprete the discontinuity at ’3’ as the signature of the endof a first substructure inside the large ICME. This substruc-ture ’2-3’ has also specific physical properties: there are nosignatures of expansion (the profile of V does not show asignificant slope), and the observed temperature decreasesin time (hotter near the sheath and colder near substructure’3-4’). As a consequence of the decreasing Tp profile, valuesof βp also decrease, being βp ∼ 10−1 near ’2’ and reachingvalues as low as βp ∼ 10−3 near ’3’. Thus, the substruc-ture ’2-3’ presents some MC signatures (low βp, coherent


rotation of the magnetic field vector) but not all ICMEsfeatures (e.g., Tp is not significantly lower than Tex). ThisMC presents a short temporal duration with a spatial sizealong the Sun-Earth direction of ∼0.09 AU (∼4 hs multi-plied by ∼900 km s−1). We interpret the short duration ofthe first MC as due to the compression made by MC2 atits trailing edge; it is expected that the expansion rate (andthe consequent cooling) of the first cloud diminished duringpart of its journey, where the second MC is pushing it frombehind [Wang et al., 2003], consistently with observations(Tp ∼ Tex).

After the MD at ’3’, θB continues increasing up to ’4’,where there is another (weaker, but significant) MD. Theregion from ’3’ to ’4’ shows the characteristics of the back offlux ropes, previously found by Dasso et al. [2006] in a sig-nificantly large non-expanding magnetic cloud (Oct. 1995)and by Dasso et al. [2007] in another huge magnetic cloud(Nov. 2004) in strong expansion. The formation of this backfeature is a consequence of previous magnetic reconnectionbetween the front of the flux rope and its environment; mag-netic flux is removed from the flux rope front (the front ispeeled) while its counterpart in the rear part still remains(Figure 6 of Dasso et al. [2006]). The formation of a back infast ICMEs is also supported by numerical simulations (seeFigure 4 of Wu et al. [2005]).

At the difference of previously studied MC-backs, the re-gion from ’3’ to ’4’ is compressed by the second MC. Thisregion of interaction does not have the same characteristicsas corotating interaction regions [CIRs, e.g. Pizzo, 1994] ortypical merged interaction regions [MIRs, e.g. Burlaga andNess, 1993, and references there in]. CIRs are due to a fastSW overtaking a slow SW, and they develop at larger dis-tance from the Sun (≥ 10 AU) than MIRs. The physicsinvolved in the region between two interacting MCs is ex-pected to be different, since MCs have a moderate spatialextension and are structured by the magnetic field. It im-plies that the compressed plasma can be evacuated on thesides, and that the second MC is able to accelerate fully thefirst one. Without a significant reconnected flux betweenthe two MCs (when the magnetic fields are nearly parallel)and after a transient period of time from the interaction, thetwo MCs are expected to travel together (a situation fullydifferent than in the case of a fast SW overtaking a slowSW).

The next substructure that can be identified is ’4-7’. Itis a very huge region (∼ 0.3 AU) that presents a low vari-ance of ~B and clear signatures of an expanding MC, witha very strong magnetic field that rotates coherently (fromnorth-east to south-west). This region shows an almost lin-ear decay of B (from B ∼ 60 to B ∼ 15 nT) consistent withthe observed expansion, with a linear V profile from 900 to700 km s−1 during ∼ 14 hours (equivalent to ∼ 0.3 AU),a typical expansion rate observed in MCs [Demoulin et al.,2008]. It also presents values of Tp significantly lower thanTex, low proton density (np ∼ 3 cm−3), and very low valuesof βp (∼ 10−3). All these signatures discard other well-known non-MCs structures (as e.g. corotating interactionregions).

Inside region ’4-7’, we also remark ’5’ and ’6’. At ’5’ astrong decrease of np starts and there are small variationsin the large scale trend of θB , φB , B, and V . At ’6’, βp, np,and Tp have peaks (here Tp reaches Tex). Thus, it is a prioriunclear where to set the end of the second flux rope, we candistinguish three possible ends: ’5’, ’6’, or ’7’. Finally, weset it at ’7’ because, at this position, a very strong MD isfound and the coherence of ~B is lost. The presence of twoflux ropes inside this ICME is also supported by applyingthe Grad-Shafranov technique [Hu and Sonnerup, 2002] tothe plasma and magnetic field data at L1 (Mostl, C., 2007,private communication).

An expanding structure with ICME signatures (e.g., lowproton temperature, low βp, decreasing V and B profiles)

still remains after ’7’. At position ’8’, there is another MD(observed mainly in θB and φB) together with a change inthe decay rate of B. This MD is associated with a sud-den increase of Tp (and consequently of βp, because np re-mains approximately constant), while the velocity profilestill presents roughly the same slope (i.e. the same expan-sion rate). Even when the expansion signatures (in both,V and B) and the coherence of ~B end at ’9’, low βp andTp < Tex remain along one hour, until ’10’. Here B recoversvalues as low as those found in the typical solar wind, Tp re-covers the expected values (Tp ∼ Tex), the coherence in therotation of ~B is lost, and the level of fluctuations (e.g., in ~B,see panels θB and φB) start to increase significantly. Thislarge region, from ’7’ to ’10’, shows the characteristics of anextended back feature belonging to the second flux rope.

The last panel of Figure 1 shows the geomagnetic effectof the ICME, as monitored by the Dst index (preliminaryindex downloaded from the OMNI data base). From ’2’ theintensity of the storm increases, because of the high intensityof the driven electric field (V Bs ∼ 5× 104 nT km s−1, withBs being the southern component of interplanetary mag-netic field).

The storm starts its recovery phase at ∼ 08:00 UT,when the magnetospheric response is the strongest (reachingDst = −263 nT); we notice a significant change in the decayrate after ’3’. Later, Dst remains < −100 nT during partof the recovery phase (until ’5’). Assuming a pure decay,decay times of τ ∼ 5 hours and τ ∼ 17 hours are found forthe first (from the peak of the storm to ’3’) and second stage(beyond ’3’), respectively. These two values are beyond thelower and close to the upper border of the typical range ob-tained for the ring current decay time, τ = 14 ± 4 hours,by Dasso et al. [2002]. Larger decay times for the recoveryphase have been associated with the presence of multiple IPstructures near Earth [Xie et al., 2006].

2.2. First Magnetic Cloud

In this section, we model the magnetic structure follow-ing the sheath (i.e., the data between ’2’ and ’3’) that wecall the first cloud (MC1).

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Figure 2. ACE observations (dots) of the magneticfield vector components in the cloud frame (see Sec-tion 2.2) for the first flux rope (observations between ’2’and ’3’). From upper to lower panels: Bx,cloud, By,cloud,and Bz,cloud. The solid lines correspond to the fittedcurves obtained using the Lundquist’s model.


To better understand the MC properties, we define a co-ordinate system linked to the cloud in which zcloud is ori-ented along the cloud axis (with Bz,cloud > 0 at the cloudaxis). Since the speed direction of a cloud is mainly alignedwith the Sun-Earth direction and is much larger than thespacecraft speed, we assume a rectilinear spacecraft trajec-tory in the cloud frame. This trajectory defines a directiond (pointing toward the Sun). Then, we define ycloud in thedirection zcloud × d and finally xcloud completes the right-handed orthonormal base (xcloud, ycloud, zcloud).

We define the axis latitude angle (θ) as the angle be-tween the cloud axis and the ecliptic plane, and the axislongitude angle (ϕ) as the one between the projection of thecloud axis on the ecliptic plane and the Earth-Sun direction(xGSE) measured counterclockwise [see Dasso et al., 2006].We also define the impact parameter, p, as the minimumdistance from the spacecraft to the cloud axis.

The local coordinate system is especially useful when pis small compared to the MC radius (R). In particular, forp = 0 and a cloud described by a cylindrical magnetic con-figuration ~B(r) = Bz(r)z + Bφ(r)φ, we have xcloud = r andycloud = φ when the spacecraft leaves the cloud. In this par-ticular case, the magnetic field data will show: Bx,cloud = 0,a large and coherent variation of By,cloud (with a change ofsign), and an intermediate and coherent variation of Bz,cloud,from low values at one cloud edge, achieving the maximumvalue at its axis, and returning to low values at the otheredge.

The minimum variance (MV) method [Sonnerup andCahill , 1967] has been used to estimate the orientationof MCs [see e.g., Bothmer and Schwenn 1998, Leppinget al. 1990, Farrugia C. J. et al. 1999, Dasso et al. 2003,Gulisano et al. 2005]. In particular, using a synthetic set ofideal cylindrical clouds, Gulisano et al. [2007] have shownthat the application of the MV technique to the normal-ized observed time series of the magnetic field ( ~B(t)/| ~B(t)|)can provide very good estimations of the cloud axis (when0 ≤ p/R ≤ 0.7). This is feasible because when ~B(t) is nor-malized, the information of the rotation of the field vector isnot mixed with possible changes in its absolute value duringthe observations.

We apply the normalized MV method to the observationsbetween ’2’ and ’3’. We find a left handed flux rope, ori-ented such that θMV =-15◦ and ϕMV =125◦, consistent withthe SEN orientation determined in Section 2.1. From thisorientation and the mean bulk speed (894 km s−1 for thisrange), we estimate the cloud size perpendicular to its axisas ∼ 0.08 AU.

In the MC frame (see Figure 2), the magnetic field hasthe typical shape observed in clouds. The almost con-stant Bx,cloud profile (with a mean value < Bx,cloud >∼−11.5 nT), indicates that p is not zero. Moreover, a roughestimation of p can achieved using the expression found byGulisano et al. [2007] (p/R ∼

√(< Bx,cloud > /B0)/1.6),

which gives p ∼ (0.3− 0.4)R.The MC borders set as ’2’ and ’3’ are in agreement with

the expected conservation of magnetic flux across a planeperpendicular to xcloud [Dasso et al., 2006], i.e., similar ar-eas below and above the curve By,cloud which equals beforeand after the MC center. These boundaries are also in agree-ment with the expected MD (current sheets) at the inter-faces between two structures with different connectivity (asthe boundaries of a flux rope forming a MC).

The MC axis, with ϕMV = 125◦, is dominantly pointingtowards yGSE with a significant contribution toward -xGSE .This is consistent with the spacecraft passing through theright (west) leg of the flux rope and, thus, with the MCapex located toward the left of L1 (observed from Earthto the Sun, with north upward). From the sign of the ob-served < Bx,cloud >, the cloud axis is toward the south ofthe ecliptic.

Fixing the orientation provided by the MV method, we fitthe free parameters of Lundquist’s model (see Section 1 for

the meaning of the free parameters) and obtain B0 = 57 nTand α = −40 AU−1. Solid lines in panels By,cloud andBz,cloud in Figure 2 show the curves obtained from the fit-ting, which are in a very good agreement with the observa-tions (dots).

To validate our previous results, we perform a simulta-neous fitting (SF) of the geometrical and physical free pa-rameters using the same procedure and numerical code asDasso et al. [2006]. We obtain θSF=-12◦, ϕSF=129◦, RSF =0.04 AU, p/RSF = 0.2, B0,SF = 59 nT, αSF = −42 AU−1.The agreement between the results of the SF and those fromthe MV method followed by the fitting of the physical cloudparameters is well within the precision of the methods (the

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Figure 3. Magnetic field vector in GSE componentsfrom ACE data (dots) for the second MC (between ’4’and ’7’) and fitted curve (solid line) using the model ofHidalgo et al. [2002b].

Figure 4. GOES lightcurve from May 13 to 14. Theflares related to the eruptions we consider to be thesources of the two MCs (see Sections 2.2 and 2.3) arepointed by arrows at peak intensity (13:04 UT and16:57 UT for the first and second flare, respectively).


differences are only ∼ 4◦ in the orientation, 5% in the radiusand 3% in B0). Due to the radial propagation of the solarwind and due to a possible meridional stratification of the so-lar wind properties, distortions from a cylindrical cross sec-tion are expected [see, e.g., Liu et al., 2006b]. From a mul-tispacecraft analysis of the first MC observed by STEREO,an oblated transverse size with the major axis perpendicu-lar to the Sun-Earth direction was found [Liu et al., 2008b].The good match between the observations and the cylindri-cal model (Figure 2) indicates that this flux rope could havean almost cylindrical configuration (a significant ellipticitywould have consequences on the magnetic field profile, seeVandas and Romashets [2003]). Therefore, we conclude thatthe main characteristics of MC1 are well determined.

2.3. Second Magnetic Cloud

In this section we model the second MC (MC2) observedfrom ’4’ to ’7’. A decreasing velocity profile (see Section 2.1)indicates that the cloud is in expansion. The sudden changeof θB at ’7’ is the most significant of all changes observedin the field after ’4’; furthermore, after this discontinuity,the expected coherence of ~B is lost. Thus, as previously dis-cussed in Section 2.1, we choose ’7’ as the rear boundary ofMC2.

The use of the MV method for this second MC does notprovide meaningful results, since the impact parameter isvery large as indicated by the low rotation of the magneticfield vector (Figure 1) and also because Bx,cloud is the largestfield component in the cloud frame [not shown]. Moreover,there is significant expansion, thus normalizing the field isnot enough to fully remove this effect. The SF to Lundquistsolution, even with a normalized field, cannot provide in thiscase a reliable result for the same reasons. When the impactparameter is low and the boundaries of the flux rope are welldetermined, flux rope modelling generally provides a goodrepresentation of the magnetic field configuration of a MC[Riley P. et al., 2004]. However, models need to be tested us-ing simultaneous observations of different parts of the sameflux rope, e.g. as recently done using STEREO observations[Liu et al., 2008b] confirming the flux rope geometry of thestudied event. Since simultaneous observations from space-craft with a significant separation are not available for thisevent, we apply a different model and method for MC2, de-scribed below. This model has more freedom in comparisonwith Lundquist’s model used before, such as the oblatenessof the cloud cross section [Liu et al., 2006b].

Observations of some expanding MCs traveling in thesolar wind are consistent with cylindrical expansion [e.g.,Nakwacki et al., 2005; Dasso et al., 2007; Nakwacki et al.,2008]. However, we can anticipate distortions from the cylin-drical shape for MC2 because of the presence of MC1 in itsfront. Therefore, we compare observations with the modelby Hidalgo et al. [2002b], which considers a magnetic fieldwith an elliptical cross section. This model has eight freeparameters: the magnetic field strength at the cloud axis,the latitude (θ) and longitude (ϕ) of the cloud axis, theimpact parameter (p), the orientation of the elliptical crosssection relative to the spacecraft path (ζ), a parameter re-lated to the eccentricity of the cross section (η), and twoother parameters related to the plasma current density.

The fitted curves using the model of Hidalgo et al. [2002b]are in very good agreement with observations (Figure 3).The model corresponds to a left handed magnetic cloud withθ = (55± 10)◦ and ϕ = (120± 10)◦.

3. Solar clues for two source events

3.1. The data

In this section, we analyze solar data from the photo-sphere to the upper corona in search of two candidate solar

events that could be the sources of the two MCs describedin previous sections.

The photospheric magnetic field evolution is analyzed us-ing observations from the Michelson Doppler Imager [MDIScherrer and et al., 1995], on board the Solar and Helio-spheric Observatory (SOHO), which measures the line ofsight magnetic field at the photosphere. These data are theaverage of 5 magnetograms with a cadence of 30 seconds.They are constructed once every 96 minutes. The error inthe flux densities per pixel in the averaged magnetograms is∼ 9 G, and each pixel has a mean area of 1.96 Mm2.

At the chromospheric level we have used full-disk Hα datafrom Big Bear Solar Observatory (BBSO) and the OpticalSolar Patrol Network (OSPAN) at the National Solar Ob-servatory in Sacramento Peak.

To identify changes in the low corona associated tothe source regions of the identified MCs, the Extreme-Ultraviolet Imaging Telescope [EIT, Delaboudiniere J.-P. etal., 1995] on board SOHO is used. EIT images chromo-spheric and coronal material through four filters. In par-ticular, we have analyzed the 195 A band which imagesplasma at 1.5×106 K. When available we have used observa-tions taken by the Transition Region and Coronal Explorer[TRACE, Handy , 1999] in the 171 A band.

The identification of CMEs is done using LASCO onboard SOHO, plus proxies for eruptions in the chromosphereand lower corona. For the analyzed time interval, LASCOimaged the solar corona from ∼ 2 - 30 solar radii with twodifferent coronagraphs: LASCO C2 (2 – 6 solar radii) andC3 (4 – 30 solar radii).

3.2. Solar activity from May 11 to May 12, 2005

Without any doubt, the main contribution to the largercloud (MC2) and other IP signatures discussed in previ-ous sections is the halo CME appearing in LASCO C2 at17:22 UT on May 13 (see Sections 3.4 and 3.5). However,since we have found that the ACE magnetic field and plasmaobservations from 05:42 UT on May 15 (’3’) to 10:30 UT onMay 17 (’10’) can be interpreted as being comprised by twodifferent structures (see Sections 2.2- 2.3), we describe thesolar activity observed by EIT and LASCO preceding themajor CME on May 13.

From May 11 until May 13, 2005, solar activity is mainlyconcentrated in two ARs, AR 10758 located in the south-ern hemisphere and AR 10759 in the northern hemisphere,which produces the most intense events. Activity progres-sively increases in the later region from May 11 to 12; then,flares reach level 2B in Hα and M1.6 in soft X-rays.

On May 11 a CME is first seen in LASCO C2 at 20:13 UTabove the SW limb. This event is classified as a full halo.The linear and second order fittings to LASCO C2 ob-servations give plane-of-sky (POS) speeds of 550 km s−1

and 495 km s−1, respectively (from LASCO CME Cata-log, http://cdaw.gsfc.nasa.gov./CME list/index.html). AnM1.1 flare in NOAA AR 10758, located at S11W51, start-ing at 19:22 UT can be associated with this event. Sub-frames from the EIT shutterless campaign show a raisingfeature in absorption by 19:22 UT. However, the closenessin time/space between the two MCs discussed in Section 2.1makes it highly improbable for this event to be the source ofMC1, since by simple assumption of constant speed it wouldtake ∼ 3.1 days to reach Earth and then would be expectedto arrive before May 15, i.e. earlier than MC1.

On May 12, an event (flare, dimming, and loops in expan-sion) is seen in EIT at about 1:57 UT. This event originatesin AR 10759 and is accompanied with a C9.4 flare locatedat N11E31 with peak at 1:10 UT. No CME is observed in


Figure 5. OSPAN Hα images of the flare at 12:49 UT. The left panel shows the portion of the filamentlying along the E-W northern magnetic inversion line expanding northward. Flare ribbons and flare loopslinking the ribbons are visible in the right panel. The E-W middle portion of the filament is no longervisible.

LASCO. However, a diffuse and gusty outward flow is ev-ident in LASCO C2. There is another apparent event inEIT at about 13:56 UT, to the south of AR 10760. The re-lated C3.0 flare occurs in S11W62 at 13:40 UT. This eventshows an association with a CME, weak and slowly travel-ing in the SW direction. This CME has a strong componentin the POS and, therefore, we do not expect it to be sig-nificantly directed towards Earth. Furthermore, from theorientation of AR 10760 polarities during their disk transit,we can infer from MDI data that its magnetic helicity signis positive, contrary to that of both studied MCs [see thedescription and interpretation of magnetic tongues in LopezFuentes et al., 2000, 2003].

After discarding these previous events, and consideringthe closeness in time/space between the two MCs analyzedin Section 2, we concentrate in the activity observed duringMay 13. In the next sections we will present observationsand discuss, in particular, two flares and associated filamenteruptions that we consider to be the solar source events ofthe clouds. The flare timings are given by GOES soft X-raydata (Figure 4).

3.3. May 13 events: Photospheric and chromosphericsignatures

Along May 13 the filament in AR 10759 is seen to activateseveral times, mainly in its northern fraction. At 12:49 UTa C1.5 flare starts in N17E15, with peak at 13:04 UT. Therising filament is seen in Hα and also in TRACE (see theleft Figure 5 and Section 3.4). At the chromospheric levela two-ribbon flare, classified as a subflare faint, developsalong the magnetic inversion line at the north of AR 10759(Figure 5). This is the only event showing eruptive char-acteristics at that time on the solar disk. This portion ofthe inversion line is oriented in the E-W direction. Part ofthe filament extending along this E-W inversion line is nolonger visible by 12:54:02 UT. Flare loops linking the tworibbons are also visible in Hα at this time. What may bethe cause of the destabilization of the E-W portion of thefilament? We have analyzed the magnetic changes observedin the filament channel previous to the flare at 12:49 UT, forthis we use MDI line of sight magnetic maps. Continuousflux cancellation is observed at the magnetic inversion linealong which the erupting part of the filament lies. The E-Wnorthern fraction of the filament shows a non-uniform shape

(as being formed by several sections) and some barbs or feetare present (Figure 6 left). A small positive polarity is seento intrude in the negative field in between two sections ofthe filament (see the arrow in Figure 6 center); this singleintrusion implies a cancellation of ∼ 5× 1019 Mx. The closerelationship between the flare and the magnetic cancellationsite is shown in the right panel of Figure 6. MDI low tempo-ral and spatial resolution does not allow us to follow in detailthe evolution of the small flux concentrations and ascertainthe clear association between their cancellation and filamenteruption; however, we believe this to be the most plausiblecause. Small-scale magnetic changes in flux concentrationsalong filament channels have been frequently reported a fewhours prior to local filament restructuring [see e.g., Chaeet al., 2000; Deng et al., 2002; Wood and Martens, 2003;Schmieder et al., 2006]. Furthermore, Zhang et al. [2001]attributed the origin of the major solar flare and filamenteruption on July 14, 2000, to flux cancellation at many sitesin the vicinity of the active region filament.

Taking into account that the C1.5 flare occurs by onlyabout 4 hours earlier than the main flare and eruption inAR 10759 (see below) and the similar orientation of theMC1 axis and of the magnetic inversion line along whichthe erupting filament lies (Figure 6 left), we consider thatthe partial eruption of the AR filament is the source of thisfirst cloud. There is evidence for the breakup of filamentsinto more than one segment, [see e.g. Martin and Ramsey ,1972] and a recent example in Maltagliati et al. [2006]. Evensome filaments erupt only partially, or saying it in a differ-ent way different portions of the same filament may erupt atdifferent times and trigger different flares [Martin and Ram-sey , 1972; Tang , 1986; Maltagliati et al., 2006; Gibson et al.,2006a; Gibson and Fan, 2006; Liu et al., 2008a]. This kindof eruptions have been discussed in the frame of the eruptiveflux rope model for CMEs [see the review by Gibson et al.,2006b]. Moreover, as discussed below in Section 3.5, on May13 there are no other evidences of CMEs in LASCO until17:22 UT.

The solar event following the one just described is themajor M8.0/2B flare that starts at 16:13 UT and peaks at16:57 UT. This flare ended on May 14 at ∼ 17:00 UT, be-ing a long duration event (see Figure 4). The photosphericand chromospheric observations corresponding to this eventhave been shown and discussed by Yurchyshyn et al. [2006]


Figure 6. Evidence of flux cancellation at the northern E-W portion of the filament channel in AR 10759.Left: BBSO Hα image early on May 13 showing the non-uniform shape of the E-W filament portion. Thetwo thick white arrows indicate the axis orientations of the two identified MCs projected on the solardisk. Center: MDI magnetogram closest in time to the flare at 12:49 UT (the field has been saturatedabove/below ± 100 G). The arrow points to a small positive flux concentration that moves towards thenegative field across the filament channel. Several of such flux cancellation episodes occur during May 13.Right: An overlay of a flare image and the ± 50 G MDI isocontour (white positive field, black negativefield). In all images the axes are in arc sec.

Figure 7. TRACE images in 171 A. The first three panels show different stages of the expansion ofthe filament embedded in a brighter curved structure during the flare at 12:49 UT. The last panel showsthat the coronal loop structure recovers its global shape after the C1.5 flare and before the M8.0 flare at16:32 UT. The same region is showed in the four panels, though the TRACE field of view is shifted forthe image at 16:08:43 UT.

Figure 8. (a) SOHO/EIT difference image in the Fe XII bandpass (195 A) showing a large coronal dim-ming (dark areas) corresponding to the CME associated with the M8.0 flare on May 13, 2005. The lastpre-eruption image taken at 16:26 UT was subtracted from the image taken at 17:07 UT. (b) SOHO/EITimage in the Fe XII bandpass (195 A) taken at 22:57 UT on May 13, 2005 (after the start of the CMEassociated with the M8.0 flare). Post-eruption arcade and a transient coronal hole are marked with topand bottom arrows respectively.

and Liu et al. [2007]. The former authors have associatedthis flare and eruption to a single MC comprising the two

we have identified in this paper. We want to stress here that


Figure 9. LASCO images of the halo CME on May 13, 2005, as recorded by a) LASCO C2 and b)LASCO C3. A pre-event image has been subtracted in both cases. These are the two only full-field ofview images taken of the event. The arrows indicate the edge of the faintest front.

the direction of MC2 axis is in agreement with the orienta-tion of the N-S fraction of the AR filament that lies alongthe main magnetic inversion line (see Figure 6, left panel).

3.4. May 13 events: The lower corona signatures

In the lower corona, TRACE observes the event at12:49 UT in 171 A (Figure 7). Images before 12:39:49 UTand after 13:04:27 UT are not useful because of a strongparticle shower. TRACE data are dominated by a large-scale sigmoidal structure (with an inverse S shape) previousto the flare well visible by ∼ 12:40 UT. The small scaleloops linking the two Hα ribbons are not distinguishable inTRACE images. By 12:49:43 UT the filament can be ob-served as a dark curved structure surrounded by a brighterone, probably formed by heated filament material (see Fig-ure 7). This bright structure is seen expanding upward until13:04:27 UT, though it is hard to assure that it erupts be-cause of lack of visibility in the following images (Figure 7,central panels).

Later on, by ∼ 16:00 UT (last panel in Figure 7) and be-fore the M8.0 flare, TRACE images appear again dominatedby the large-scale sigmoid. This structure is similar to theone described by Liu et al. [2007] seen before the M8.0 flare;in their Figure 1 they show the different parts that consti-tute the sigmoid: the magnetic elbows, the envelope loops[according to the nomenclature of Moore et al., 2001] andthe footpoints of the sigmoidal fields. Thus, the rising ofthe bright curved structure, shown in the central panels ofFigure 7, is most likely linked to the filament eruption duringthe first and less intense two-ribbon flare that we associatewith MC1. Yurchyshyn et al. [2006] have suggested that thebright structure seen in the central panels is the core of thesigmoid that later starts a slow ascent (by ∼12:55 UT) andbecome active during the M8.0 flare, but in fact it eruptedearlier and reactivated later.

The major eruptive event, linked to an M8.0 flare, is visi-ble in EIT images starting from 16:37 UT (see Figure 8). Itexhibits several CME-associated phenomena [Hudson andCliver , 2001]: coronal dimmings, a large-scale EIT waveand a post-eruption arcade. As it was pointed out by Liuet al. [2007], the darkest part of the dimming comprises twopatches to the NE and SW of the erupting active region.The SW dimming is associated with a transient coronal hole(shown by an arrow in Figure 8b). It was suggested thattransient coronal holes can be interpreted as footpoints ofthe ejected flux rope [see, e.g., Webb et al., 2000; Mandriniet al., 2005]. The absence of the clearly visible second tran-sient coronal hole in this event may indicate that the secondleg of the erupted flux rope may have already reconnected

with the ambient coronal magnetic field [this reconnectionchanges the very extended field lines rooted in the dimmingby closed magnetic connections Attrill et al., 2006; Kahlerand Hudson, 2001]. The overall morphology of twin dim-mings and the presence of the post-eruption arcade (Fig-ure 8b) suggest that an eruption of a coronal magnetic fluxrope had taken place [e.g., Hudson and Cliver , 2001].

3.5. May 13 events: The upper corona observations

On May 13, 2005, a full halo CME with brightness asym-metry towards the NE first appears in the field of view ofLASCO C2 at 17:22 UT (see Figure 9a). Five more partialsnapshots of the halo CME were captured by LASCO C3from 17:20 to 17:50 UT. The closeness in time of the lat-ter images does not allow to see major changes from one toanother, so we present here only one of them (Figure 9b).The estimated POS speed for the fastest front, as calculatedby the LASCO CME catalog, is ∼1690 km s−1. Further in-formation on the kinematics of this CME and its associatedshock can be found in Section 4 and in Reiner et al. [2007].

Enhancement of the LASCO images presented in Figure 9allows to distinguish a very faint edge toward the south andnorth, about a couple of solar radii ahead of the bright-est one (indicated by arrows). This feature suggests thepresence of a shock, as described by Vourlidas et al. [2003]and Vourlidas [2006]. Finally, the CME related to the solareruption at 12:49 UT is probably so faint, that its detectionwas not likely due to sensitivity limitations of LASCO [e.g.,Tripathi et al., 2004; Yashiro et al., 2005]. Furthermore,Yashiro et al. [2005] conclude that all fast and wide CMEsare detectable by LASCO, but slow and narrow CMEs maynot be visible when they originate from disk center. In ad-dition to this restriction, which impedes the visualization ofa CME if it is feeble, the LASCO cadence during and afterthe C1.5 flare at 12:49 UT was of 30 minutes, as opposite tothe usual 12 minutes of LASCO C2.

4. Two structures in the inner heliosphere:Remote radio type II observations

When a halo CME leaves the field of view of a corona-graph, it is occasionally possible to further track its evo-lution by means of measurements in the radio regime ofthe electromagnetic spectrum. The shock commonly drivenby a CME excites electromagnetic waves emitted at the lo-cal plasma frequency (f ∼ √

ne) and/or its harmonic. As


Figure 10. Dynamic spectral plot showing the radio emission recorded by the WAVES/TNR receiveron board Wind between May 13 at 12:00 UT and May 15 at 12:00 UT. The vertical axis represents 1/f[kHz−1] while the horizontal axis corresponds to time. Three identified type II emissions are labeled (theentity (a) is associated to (1) and (2) which are the fundamental and harmonic emission at the plasmafrequency, and the entity (b) is associated to (3)). The black arrow in the horizontal axis denotes theshock arrival time. Tick marks on the right vertical axis indicate the location of TNR frequency channelsfor the plotted frequency range.

the shock travels outward from the solar corona, the localplasma frequency decreases, and a drifting signal, a type IIradio burst, is generated. This brand of emissions can befollowed drifting in frequency from ∼ 150 MHz down to∼ 25 kHz, which is the approximate value of the local plasmafrequency in the vicinity of Earth [Leblanc et al., 1998].Then, the shock traveling toward Earth is detected by in-struments in the IP space, such as WAVES on board Wind[Bougeret J. L. et al., 1995]. The ICME can also be de-tected by ground-based antennas measuring interplanetaryscintillation of radio sources [e.g., Manoharan, 2006].

The metric components of type II bursts have long beenused by space weather forecasters as input to models ofshock arrival time prediction [e.g., Dryer and Smart , 1984;Smith and Dryer , 1995; Fry et al., 2001]. Nevertheless, thereis ongoing debate on the mechanism that generates them,questioning their direct relationship with shocks ahead ofCMEs [e.g., Pick , 1999; Gopalswamy et al., 2001a; Claßenand Aurass, 2002; Cane and Erickson, 2005]. Therefore,Cremades et al. [2007] attempted to improve space weatherforecasts by employing information of type II radio burstsin the kilometric domain (30 − 300 kHz), which are indeedbelieved to be generated by IP shocks [Cane et al., 1987].Accordingly, the slope of the drifting radio emission can betaken as a proxy for the associated shock speed [Reineret al., 1998]. Furthermore, for each point of emission de-tected in the space-borne radio receivers, it is possible toobtain a proxy of the radial distance from the Sun throughdensity models [e.g., Saito et al., 1977; Leblanc et al., 1998;Hoang et al., 2007]. The technique used to characterize thetemporally-associated radio signals makes use of the densitymodel developed by Leblanc et al. [1998]. To obtain the ap-proximate sites of the emissions, and consequently a proxyof the shock speed, it is only needed to enter into the densitymodel an estimate of the plasma density near Earth, usu-ally between 5 and 9 cm−3. The top frequency taken intoaccount as input for the technique (300 kHz for the funda-mental emission) corresponds approximately to a distance of20 solar radii according to that heliospheric density model.

We assume that at those heights most of the decelera-tion has already taken place, and thus that the slope of thelinear profile in Figure 10 is enough to characterize the kine-matics of a shock that travels at a constant speed. Beyond

20 solar radii, the corresponding frequency channels of theWAVES/RAD1 detector are scarce in comparison with thoseof WAVES/TNR (see rationale at the end of this Section).That is why the main goal of the latter detector is the track-ing of radio phenomena up to the orbit of Earth. Furtherdetails on the technique, its application and success relativeto other methods can be found in Cremades et al. [2007].

After the solar events on May 13, a number of type IIradio features in the kilometric domain were observed bythe Thermal Noise Receiver (TNR) experiment on WAVES.Figure 10 displays the radio intensity between May 13 at12:00 UT and May 15 at 12:00 UT, and from f ∼250 kHz tof=20 kHz. In addition, each of the identified low-frequencytype II emissions has been labeled as follows: (1) a radioemission drifting at a speed of 1400 km s−1, (2) an emissiondrifting at a speed of ∼1450 km s−1 and certainly being theharmonic of the first feature, and (3) a slower radio emis-sion drifting at a speed of ∼850 km s−1. The application ofthe technique described in Cremades et al. [2007] to emis-sion (1) yielded for its associated shock a predicted arrivaltime at ∼02:40 UT on May 15, which is in good agreementwith the real shock arrival at Earth (02:11 UT). The arrivaltime derived from radio signal (2) yielded ∼23:20 UT onMay 14. The error of ∼3 hours might be attributed to the’patchiness’ of the emission. This effect arises because a sin-gle shock may drive type II bursts at multiple places alongthe shock surface, with the possibility of a large non-radialcomponent of the shock velocity. In addition, there could besignificant differences in the drift rates if the density gradi-ent or the local shock speed varies along the shock surface.Under these circumstances entities (1, 2) and (3) could bethe manifestation of the shock driven by only one ICME.However, we have no data to justify any of these hypoth-esis. Moreover, the large difference in propagation speedbetween entities (1, 2) and (3) suggests that they must beindeed produced by two distinct structures traveling in thesolar wind or by the same structure traveling with a sig-nificantly slower speed at a first stage of its journey andwith a faster speed at its second stage. We will interpretthis second possibility in the frame of the global observa-tions in next section (Section 5). Finally, the prediction of


the arrival time of entity (3) yielded ∼10:20 UT on May 15.The linear backtracing of this last emission in time matcheswell with the earlier solar event on May 13. On the otherhand, entities (1) and (2) show good temporal associationwith the occurrence of the second solar event on May 13 (at16:57 UT).

Based on the observations of WAVES/TNR we presenthere a different interpretation of the sources of the radiosignatures than Reiner et al. [2007], who used radio data ac-quired by the WAVES/RAD1 detector (see next paragraph).The TNR observations indicate the existence of 2 interplan-etary shock entities, which from now on will be referred toas entity (a) -the shock traveling at ∼1400 km s−1, corre-sponding to (1) and (2)- and entity (b) -the shock travelingat 850 km s−1, related to (3). Entity (b) presumably drawsnear entity (a) around 07:00 UT on May 14, at an approxi-mate distance of 98 solar radii from the Sun. There are novisible signs of ’cannibalism’ [Gopalswamy et al., 2001b] inradio data, likely due to the weakness of entity (b) and/orbecause there is no significant reconnection between the twoentities. Moreover, it is worth to note that entity (a) wasnot visible before 06:00 UT on May 14, as opposite to theslower entity (b). A plausible reason for this effect is thatentity (a) may have been moving in an environment plasmafaster than the typical solar wind, therefore without leadingshock formation.

The kinematics analysis performed for the radio event onMay 13-15, 2005 by Reiner et al. [2007] takes into accountentity (a) only. Those authors present spectral plots of theevent that reach down in frequency up to the coverage of theWAVES/RAD1 detector. As seen in Figure 10 of that study,entity (b) is not obvious from RAD1 measurements. How-ever, it does show up in TNR data, suggesting that Reineret al. [2007] did not consider those data in their analysis. Itis most likely that the poor frequency resolution of RAD1below 256 kHz is guilty of overseeing entity (b). The RAD1detector employs at a time only 32 frequency channels tomonitor the range 20-1040 kHz, using only 21 channels be-low 256 kHz and interpolating to obtain the rest of the data.This may lead to false apparent broadband emissions and tothe oversight of some narrow-band features in the RAD1 do-main. Conversely, the TNR receiver employs 96 channels tocover the frequency range of 4-256 kHz, over five logarith-mically spaced frequency bands [Bougeret J. L. et al., 1995,also see location of frequency channels as tick marks on theright vertical axis of Figure 10]. This configuration achievesbetter frequency resolution at distances greater than ∼20solar radii, and hence makes it best suitable to study emis-sions beyond those distances.

Figure 11. Schematic global view of the two MCs andthe trajectory of the spacecraft that observed them. Thefirst MC is drawn with a continuous line (MC1), whilethe second one with a dashed line (MC2).

5. Conclusion

We propose here a fully consistent physical scenario forthe chain of events from May 13 to 17, 2005, which is basedon a detailed analysis of the observations presented in pre-vious sections.

Two solar events occurred on May 13 with a differenceof ∼ 4 hours, both from AR 10759. The first one, a C1.5two-ribbon flare that was associated with the ejection of apart of the AR filament extending along the E-W portion ofthe inversion line (north of the AR). The second and largestone, an M8.0 two-ribbon associated with the ejection of thefilament lying along the N-S portion of the inversion lineand a CME observed in LASCO C2.

In the IP medium we identify an anomalously long ICME(more than 2 days). Following the results of Burlaga et al.[2001, 2003], this suggest the presence of two flux ropes (seeSect. 2.1). Indeed, we have found two clouds (MC1 andMC2). They are stacked together and only the latitude, θB ,of the magnetic field permits to identify in the data twocoherent structures separated by a region (called ’3-4’ inFig. 1) that have a different temporal evolution from neigh-boring regions. Such conclusion is supported by the variousattempts done to fit the data by one or two flux ropes. In-deed, only the two flux rope models presented in Section 2give a reasonably good fit to the data. Figure 11 shows asketch of the two clouds near Earth, in which we have takeninto account their orientations, relative sizes, and the space-craft trajectory during the observations. The two MCs areoriented such that: θMC1 = (−12± 5)◦, ϕMC1 = (125± 5)◦,θMC2 = (55± 10)◦, and ϕMC2 = (120± 10)◦. MC1 presentsa low impact parameter (p/R ∼ 0.2 − 0.3) and a radius of∼ 0.04 AU, while MC2 is much larger and presents a largerimpact parameter; both MCs have negative magnetic helic-ity. MC1 is small, has a very intense magnetic field, thesignature of an overtaking flow penetrating deep in the outbound, and almost no portion with low proton temperature(see Figure 1); all these observations can be interpreted asthe consequence of a compression by MC2.

The left panel in Figure 6 shows the projection on thesolar surface of the axes of MC1 and MC2 and the senseof their azimuthal field component, for comparison with the

12:00 00:00 12:00 00:00 12:000

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Figure 12. Cartoon representing the position (as dis-tance from the Sun along the Sun-Earth line) of theevents ejected from the Sun in function of time. Ver-tical dotted lines (ticks at S1 and S2) represent the timeof beginning of solar flares associated with both ejections(May 13, ∼ 13:00 UT and ∼ 16:00 UT, respectivelly.Vertical dashed lines (ticks at Shock, MC1, and MC2)represent the arrival of the shock (May 15, 02:11 UT)and of the center of magnetic clouds 1 and 2 (May 15, ∼09:00 UT and ∼ 21:00 UT, respectively).


orientation of the erupting parts of the filament associatedwith the C1.5 and M8.0 flares, respectively. Only the sec-ond flare has an observed associated CME, and it is apriorisurprising to associate a C1.5 flare to a fast ejection (witha velocity ≈ 850 km s−1 from radio data); however, it hasbeen shown by Gopalswamy et al. [2003, 2005] that the cor-relation between CME speed and flare intensity is weak.Apart these two issues, the other characteristics of the solarand interplanetary events are in agreement as follows. Theerupting filament orientations and directions of the field intheir associated arcades (see right panel of Figure 5 and Fig-ure 8b) are in good agreement with the axial and azimuthalfield components in the clouds. The AR field, as well as thetwo MCs have negative magnetic helicity. Altogether, withthe transit time in the expected time interval, this demon-strates the associations between the two solar eruptions tothe two interacting MCs detected near Earth.

Moreover, an evidence of the interaction is directly foundin the interplanetary radio data, as follow. Radio type IIobservations (Section 4, entity (b)) are consistent with aspeed for the shock driven by the first ejecta (MC1) of∼ 850 km s−1, when it travels from the Sun up to ∼ 0.5 AU(from ∼13:00 UT on May 13 to ∼07:00 UT on May 14). Theshock driven by the second ejecta (MC2) seems to be weakerand is not detected using radio data in this range of time,probably because MC2 travels in a non-typical solar windmedium, i.e., a medium perturbed by the previous passageof MC1, so having a faster speed. It is also possible that thissecond shock starts to interact with the trailing part of MC1at earlier times just after its eruption and, thus, it turns tobe weaker, as shown from numerical simulations of interac-tion of flux ropes [Xiong et al., 2007]. Then, at times whenMC2 reaches MC1 (at ∼ 0.5 AU, on May 14 at ∼07:00 UT),when the exchange of momentum between the two clouds ismost efficient, MC1 is accelerated and MC2 is decelerated.Then, both start to travel at similar speeds ∼ 1400 km s−1,but keeping their individuality. The shock wave driven bythis new combined structure is the cause of the radio emis-sion observed as entity (a), identified in Section 4. Duringthe second half of the transit to Earth, MC1 is compressedfrom behind by MC2.

Summarizing, we propose the following interpretation ofthe sequence of events that occurred on May 13, 2005: twoejective solar events occurred, which could be tracked to theEarth environment and were observed as two attached, butnon-merged, magnetic clouds. Figure 12 shows the positionversus time (along the Sun-Earth line) of both interplane-tary events, as inferred from the sequence of the observationsused for our analyses from the Sun to 1 AU. The presenceof structures preserving multiple flux ropes, as the multi-ple structure studied here, has been observed in some cases[Wang et al., 2002]. Signatures of the interaction betweenejecta, similar to those found in this study, have also beenidentified from observations with the following properties:(i) acceleration of the leading ejecta and deceleration of thetrailing one, (ii) compressed field and plasma in the lead-ing ejecta, (iii) weak shock or disappearance of the shockdriven by the second ejecta, and (iv) strengthening of theshock driven by the leading one [Farrugia and Berdichevsky ,2004].

This study especially illustrates the need of combining so-lar, in situ and remote sensing of interplanetary propagationof the ejecta, in order to reveal the physical processes. Pre-vious studies of the same event, but mostly limited to one ofthese three domains have concluded on the presence of onlyone ejecta (with varying characteristics depending on thestudy). In the present case, the peculiar characteristics ofthe in situ observed magnetic field first alert us on the pos-sible presence of two MCs. We confirm this by performing adetailed modeling of the magnetic field exploring various fluxrope boundaries. This stimulates a deeper study of data in

the two other domains, with the new perspective of testingtwo ejecta. The solar data were crucial to first confirm thepresence of two eruptions, and to define the timing and thespatial organization of the erupting magnetic configurations.This brought us to carry a deeper analysis of interplanetaryradio observations. Indeed, two type II bursts, travelingat different velocities, were found. These observations alsopermit to localize where the interaction occurred (about halfway between the Sun and the Earth). Then, all the piecesof the puzzle finally fit together, re-enforcing our interpreta-tion separately in each domain. However, the possibility ofa single complex event (as proposed by [Yurchyshyn et al.,2006; Reiner et al., 2007]) cannot be fully discarded and ourproposed scenario, based on compelling evidence, must befurther investigated.

Such studies of interacting ejecta will greatly benefit fromobservations taken with spacecraft in quadrature, as alreadydone for some events [e.g., Rust et al., 2005], and as availablewith STEREO spacecrafts. This has already be done with aspacecraft configuration far from optimum conditions [Har-rison R. A. et al., 2008]. On top of providing in situ mea-surements of physical parameters at three locations (com-bining ACE with STEREO) and monitoring type II bursts,there are imminent new possibilities of imaging ICMEs intheir journey through interplanetary space. Then, we willhave the opportunity to derive precise constraints on thephysical mechanisms of interaction of an ejecta with a fasterone, or simply with a fast overtaking solar wind.

Acknowledgments. The authors would like to thank to theInternational Space Science Institute (ISSI, Bern, Switzerland)for supporting the project ’The Stages of Sun-Earth Connec-tion’, lead by C. Cid. C.C, Y.C, and E.S. acknowledge finan-cial support from the Comision Interministerial de Ciencia y Tec-nologa(CICYT) of Spain (ESP 2005-07290-C02-01 and ESP 2006-08459). S.D. thanks the Argentinean grant UBACyT X425. S.D.and C.H.M. thank the Argentinean grants: UBACyT X329 andPICT 03-33370 (ANPCyT). H.C., S.D., and C.H.M. are membersof the Carrera del Investigador Cientıfico, CONICET. B.S., P.D.,and S.P. are doing partly this project in the frame of the EuropeanNetwork “SOLAIRE”. C.H.M., B.S. and P.D. acknowledge finan-cial support from CNRS (France) and CONICET (Argentina)through their cooperative science program (No 20326). L.R. andA.N.Z. acknowledge support from the Belgian Federal SciencePolicy Office through the ESA-PRODEX programme. The workat the University of Barcelona was supported by the Ministeriode Educacion y Ciencia under the projects AYA2004-0322 andAYA2007-60724. We are grateful to the ACE/SWEPAM andACE/MAG teams, for the data used for this work. We wouldalso like to thank OmniWeb from where we downloaded the pre-liminar Dst index. The CME catalog is generated and maintainedat the CDAW Data Center by NASA and The Catholic Universityof America in cooperation with the Naval Research Laboratory.We acknowledge the use of Wind WAVES/TNR radio data. TheOptical Solar PAtrol Network (OSPAN) project, previously calledISOON, is a collaboration between the Air Force Research Labo-ratory Space Vehicles Directorate and the National Solar Obser-vatory. We acknowledge the use of TRACE data. EIT, LASCOand MDI data are a courtesy of SOHO/EIT, SOHO/LASCO andSOHO/MDI consortia. SOHO is a project of international coop-eration between ESA and NASA. The authors thank S. Hoang forproviding clear informations on Wind/WAVES and the refereesfor their helpful suggestions.

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