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Title
From the Sun to the Earth: The 13 May 2005 Coronal Mass Ejection
Permalink
https://escholarship.org/uc/item/41x2g3rm
Journal
Solar Physics: A Journal for Solar and Solar-Stellar Research and the
Study of Solar Terrestrial Physics, 265(1)
ISSN
1573-093X
Authors
Bisi, M. M.
Breen, A. R.
Jackson, B. V.
et al.
Publication Date
2010-08-01
DOI
10.1007/s11207-010-9602-8
Peer reviewed
eScholarship.org
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University of California
Solar Phys (2010) 265: 49–127
DOI 10.1007/s11207-010-9602-8
R E M OT E S E N S I N G O F T H E I N N E R H E L I O S P H E R E
From the Sun to the Earth: The 13 May 2005 Coronal
Mass Ejection
M.M. Bisi · A.R. Breen · B.V. Jackson · R.A. Fallows · A.P. Walsh · Z. Mikić · P. Riley ·
C.J. Owen · A. Gonzalez-Esparza · E. Aguilar-Rodriguez · H. Morgan · E.A. Jensen ·
A.G. Wood · M.J. Owens · M. Tokumaru · P.K. Manoharan · I.V. Chashei ·
A.S. Giunta · J.A. Linker · V.I. Shishov · S.A. Tyul’bashev · G. Agalya ·
S.K. Glubokova · M.S. Hamilton · K. Fujiki · P.P. Hick · J.M. Clover · B. Pintér
Received: 9 December 2009 / Accepted: 30 June 2010 / Published online: 3 August 2010
© The Author(s) 2010. This article is published with open access at Springerlink.com
Remote Sensing of the Inner Heliosphere
Guest Editors: M.M. Bisi and A.R. Breen
M.M. Bisi () · B.V. Jackson · M.S. Hamilton · P.P. Hick · J.M. Clover
Center for Astrophysics and Space Sciences, University of California, San Diego, 9500 Gilman
Drive #0424, La Jolla, CA 92093-0424, USA
e-mail: Mario.Bisi@aber.ac.uk
B.V. Jackson
e-mail: bvjackson@ucsd.edu
M.S. Hamilton
e-mail: sarhas@gmail.com
P.P. Hick
e-mail: pphick@ucsd.edu
J.M. Clover
e-mail: jclover@ucsd.edu
M.M. Bisi · A.R. Breen · R.A. Fallows · A.G. Wood · M.J. Owens · B. Pintér
Institute of Mathematics and Physics, Aberystwyth University, Penglais Campus, Aberystwyth,
Ceredigion, SY23 3BZ, Wales, UK
A.R. Breen
e-mail: azb@aber.ac.uk
R.A. Fallows
e-mail: raf@aber.ac.uk
A.G. Wood
e-mail: aow@aber.ac.uk
B. Pintér
e-mail: b.pinter@aber.ac.uk
M.J. Owens
Space Environment Physics Group, Department of Meteorology, University of Reading, Earley Gate,
PO Box 243, Reading, RG6 6BB, England, UK
e-mail: m.j.owens@reading.ac.uk
50
M.M. Bisi et al.
A.P. Walsh · C.J. Owen
Mullard Space Science Laboratory, University College London, Holmbury St. Mary, Dorking, Surrey,
RH5 6NT, England, UK
A.P. Walsh
e-mail: apw@mssl.ucl.ac.uk
C.J. Owen
e-mail: cjo@mssl.ucl.ac.uk
Z. Mikić · P. Riley · J.A. Linker
Predictive Science, Inc., 9990 Mesa Rim Road Suite 170, San Diego, CA 92121, USA
Z. Mikić
e-mail: mikicz@predsci.com
P. Riley
e-mail: pete@predsci.com
J.A. Linker
e-mail: linkerj@predsci.com
A. Gonzalez-Esparza · E. Aguilar-Rodriguez
MEXART, Instituto de Geofisica, Unidad Michoacán, Universidad Nacional Autonoma de Mexico,
Tzintzuntzan 310, Col. Vista Bella, Morelia, Michoacán 58098, Mexico
A. Gonzalez-Esparza
e-mail: americo@geofisica.unam.mx
E. Aguilar-Rodriguez
e-mail: ernesto@geofisica.unam.mx
H. Morgan
Institute for Astronomy, University of Hawaii, 2680 Woodlawn Drive, Honolulu, HI 96822, USA
e-mail: hmorgan@IfA.Hawaii.Edu
E.A. Jensen
ACS Consulting, 40 FM 1960 W PMB 370 Houston, TX 77090, USA
e-mail: eaj@acs-consulting.com
E.A. Jensen
MMT Observatory, Mt. Hopkins, Amado, AZ 85645, USA
M. Tokumaru · K. Fujiki
Solar-Terrestrial Environment Laboratory (STELab), Nagoya University, Furo-cho, Chikusa-ku,
Nagoya 464-8601, Japan
M. Tokumaru
e-mail: tokumaru@stelab.nagoya-u.ac.jp
K. Fujiki
e-mail: fujiki@stelab.nagoya-u.ac.jp
P.K. Manoharan · G. Agalya
Radio Astronomy Centre, National Centre for Radio Astrophysics, Tata Institute of Fundamental
Research, Udhagamandalam (Ooty), 643 001, India
P.K. Manoharan
e-mail: mano@ncra.tifr.res.in
The 13 May 2005 Coronal Mass Ejection
51
Abstract We report the results of a multi-instrument, multi-technique, coordinated study
of the solar eruptive event of 13 May 2005. We discuss the resultant Earth-directed (halo)
coronal mass ejection (CME), and the effects on the terrestrial space environment and upper
Earth atmosphere. The interplanetary CME (ICME) impacted the Earth’s magnetosphere
and caused the most-intense geomagnetic storm of 2005 with a Disturbed Storm Time (Dst)
index reaching −263 nT at its peak. The terrestrial environment responded to the storm on
a global scale. We have combined observations and measurements from coronal and interplanetary remote-sensing instruments, interplanetary and near-Earth in-situ measurements,
remote-sensing observations and in-situ measurements of the terrestrial magnetosphere and
ionosphere, along with coronal and heliospheric modelling. These analyses are used to trace
the origin, development, propagation, terrestrial impact, and subsequent consequences of
this event to obtain the most comprehensive view of a geo-effective solar eruption to date.
This particular event is also part of a NASA-sponsored Living With a Star (LWS) study
and an on-going US NSF-sponsored Solar, Heliospheric, and INterplanetary Environment
(SHINE) community investigation.
Keywords Active regions, magnetic fields · Corona, models · Corona, radio emission ·
Coronal mass ejections, initiation and propagation · Coronal mass ejections,
interplanetary · Magnetic fields, interplanetary · Magnetic fields, models · Radio bursts,
association with flares · Radio scintillation · Solar wind, disturbances · Sunspots, magnetic
fields · Velocity fields, solar wind · X-ray bursts, association with flares
G. Agalya
e-mail: gkagali@gmail.com
I.V. Chashei · V.I. Shishov · S.A. Tyul’bashev · S.K. Glubokova
Pushchino Radio Astronomy Observatory, Astrospace Center, Lebedev Physical Institute, Pushchino,
Moscow region, 142290, Russia
I.V. Chashei
e-mail: chashey@prao.ru
V.I. Shishov
e-mail: shishov@prao.ru
S.A. Tyul’bashev
e-mail: serg@prao.ru
S.K. Glubokova
e-mail: glubokovask@yandex.ru
A.S. Giunta
Department of Physics, University of Strathclyde, John Anderson Building, 107 Rottenrow, Glasgow,
G4 0NG, Scotland, UK
A.S. Giunta
Space Science and Technology Department, Rutherford Appleton Laboratory, Didcot, Oxfordshire,
OX11 0QX, England, UK
e-mail: alessandra.giunta@stfc.ac.uk
S.K. Glubokova
Pushchino State University, Pushchino, Moscow region, 142290, Russia
P.P. Hick
San Diego Supercomputer Center, University of California, San Diego, 9500 Gilman Drive #0505,
La Jolla, CA 92093-0505, USA
52
M.M. Bisi et al.
1. Introduction
The phenomenon of coronal mass ejections (CMEs), their associated interplanetary counterparts, and the effects they cause on the Earth environment are well studied but still not
completely understood (see, e.g., Gosling, 1993; Harrison, 2003; Schwenn et al., 2005;
Zhang et al., 2007a). When a CME erupts from the Sun it removes a large amount of
mass (solar plasma) and magnetic energy, and hurls it out into the interplanetary medium.
When these interplanetary counterparts are detected, they are generally deemed interplanetary coronal mass ejections (ICMEs) (Lindsay et al., 1999). A subset of ICMEs containing one or more defined flux ropes can be classified as Magnetic Clouds (MCs). This special case of flux rope within an ICME was defined by L.F. Burlaga (Burlaga et al., 1981;
Burlaga, Lepping, and Jones, 1990; Burlaga, 1995) as having strong magnetic fields (when
compared with the surroundings) displaying a large and coherent rotation, and depressed
ion temperature; thus these were named as MCs.
It is important to remember the definition of a CME (see Schwenn et al., 2005; and
references therein). We can define a CME as an observable change in coronal structure
which occurs on time scales of a few minutes to several hours and where outward motion
can be observed (Schwenn et al., 2005). It also involves the instance of a newly-appearing,
bright, white-light feature, which is seen as being discrete in coronagraph imagery. Note that
this precise definition does not specify anything of the origin/source of the “feature” nor its
nature; be it ejecta themselves or any subsequent effects driven by them (Schwenn et al.,
2005).
CMEs, particularly the larger ones, most-commonly originate from active regions (ARs)
and can leave behind signatures of coronal dimming (see, e.g., Bewsher, Harrison, and
Brown, 2008). These signatures are generally best observed in extreme-ultraviolet (EUV)
frequencies or sometimes in X-rays. CMEs can be accompanied by large flares, or indeed
initiated by flare eruptions (see, e.g., Lippiello, de Arcangelis, and Godano, 2008; and references therein). Such flares often have associated coronal radio bursts, usually of Type II
or Type III, or sometimes both (see, e.g., Gopalswamy et al., 2005). The characteristic detection of a CME while looking at the Sun is by a localised increase in the brightness of
white-light emission from the Sun due to dense features scattering more of the photospheric
white light. CMEs in general contribute to about 5% of the total mass flux of the solar wind
(see, e.g., Lang, 2001).
At around 1 AU (≈ 215.5 R⊙ ) distance from the Sun, there are various features which can
be picked up by spacecraft in-situ measurements during the passage of an ICME. The larger
ones tend to result in density increases. There is often a change in velocity; this change
is relatively sharp when the ICME velocity is much different from the “ambient” solar
wind at the time. Such changes often cause interplanetary (IP) shocks to be seen before
the passing of the ICME, and sometimes behind the ICME as a reverse shock. In addition, a sheath (a density increase usually composed of piled-up solar-wind plasma) is often
seen behind the IP shock but before the core of material passes over the spacecraft (there is
sometimes another increase in density trailing or inside the ejecta which could be associated
with an erupting filament accompanying the CME), depending on the overall orientation
and placement within the ICME structure which the spacecraft measures. An MC is defined (as previously stated) as a flux rope, but not all ICMEs have to contain a single MC;
it is possible for multiple MCs to be contained within the single passage of a “merged”
ICME structure at 1 AU distance (Burlaga et al., 1981; Burlaga, Lepping, and Jones, 1990;
Burlaga, 1995).
The response of the Earth’s space environment and upper atmosphere (the terrestrial magnetosphere, ionosphere, and thermosphere) to solar-wind impact depends critically upon the
The 13 May 2005 Coronal Mass Ejection
53
orientation of the interplanetary magnetic field (IMF) carried by the solar wind (see, e.g.,
Brekke, 1997). When the north – south component of the IMF (Bz ) is negative (that is, the
field is oriented southwards), then direct reconnection can take place between the interplanetary and geomagnetic fields along the sub-solar region of the solar-wind – magnetosphere
interface – the magnetopause (see, e.g., Reiff and Burch, 1985; Schwenn et al., 2005).
Schwenn et al. (2005) also note that since the 1970s there have been two fundamentally
different recognised sources of negative Bz : sources on the non-active Sun (e.g. high-speed
streams), and sources on the active Sun (e.g. CMEs). The ionospheric footprint of merged
flux tubes are then dragged across the polar regions, driving dayside – nightside convection
in the high-latitude magnetosphere and ionospheric polar caps, with return flow at lower latitudes (Dungey, 1961). The solar wind then continues to drag the merged flux tubes downstream away from the Earth, forming a long magnetotail on the anti-solar side. We define
“downstream” in the direction of flow of the solar wind, and “upstream” in the direction
against the flow of the solar wind, i.e. back towards the Sun, and stick to this convention
throughout all aspects of this paper. Reconnection eventually takes place in the magnetotail,
separating the interplanetary and geomagnetic fields, and the resulting contraction of the
nightside geomagnetic flux tubes accelerates plasma back towards the Earth, depositing energy in the upper atmosphere around the nightside auroral oval and perturbing ionospheric
and thermospheric structure (see, e.g., Brekke, 1997). This is the situation under which the
most-efficient coupling between the solar wind and the Earth’s space environment takes
place, and leads to the most-pronounced observable effects. The precise morphology of the
magnetospheric and ionospheric convection patterns also depends on the azimuthal component (By ) of the IMF, with positive azimuthal values shifting the entire convection pattern
towards the dawn section, and negative values leading to a duskwards bias.
When the north – south component of the IMF is northward, by contrast, the interplanetary and geomagnetic fields do not interact with each other at the leading edge of the magnetosphere, but some reconnection may still occur around the polar lobes of the geomagnetic
field if the azimuthal component of the IMF is not too large (see, e.g., Freeman et al., 1993),
and under appropriate conditions, plasma of dayside origin can still be observed in the nightside ionosphere (Wood et al., 2008). Under these circumstances, the resulting convection
patterns in the magnetosphere and high-latitude ionosphere display sunwards flow around
the sub-solar point (Lockwood and Moen, 1999). Also under these conditions, energy deposition into the upper atmosphere takes place on the sunward side and energy transfer from
the solar wind to the terrestrial space environment is much less efficient: auroral activity can
still take place under conditions of northward IMF, but it is both weaker and considerably
harder to observe (see, e.g., Brekke, 1997).
Once magnetic reconnection has occurred downstream from the Earth, the magnetic field
is convected away from this region. The field on the earthward side of the reconnection point
is transported towards the Earth and plasma is accelerated in this direction. As the plasma
travels earthward, it encounters a steady increase in the strength of the geomagnetic field
and this gradient is approximately perpendicular to the direction of the field. The resulting
force on the charged particles causes protons to drift westwards and electrons to drift eastwards, setting up a current in the magnetosphere flowing from east to west known as the
ring current. This current, generally located between 4 Earth radii (R⊕ ) and 6 R⊕ , creates a
magnetic field which opposes the geomagnetic field and, following a magnetic storm, this
depression may last for several days (Hargreaves, 1992).
The response of the terrestrial ionosphere to large-scale disturbances in the solar wind
has been reviewed in detail by Buonsanto (1999) and Mendillo (2006). In summary, compression of the magnetosphere by the solar wind causes downwelling of the ionospheric
54
M.M. Bisi et al.
plasma. At lower altitudes plasma loss by recombination is enhanced due to the greater thermospheric densities and so the plasma density is reduced. The thermospheric composition
is also modified by the neutral winds and therefore the response of the ionosphere varies
with location. In addition, at polar latitudes, the IMF carried by the solar wind is connected
to the geomagnetic field and the convection electric field drives the high-latitude plasma
convection pattern. During storms this electric field intensifies, enhancing plasma transport,
and the convection pattern also expands to lower latitudes.
We report on a series of remote-sensing observations and in-situ measurements of the 13
May 2005 CME, its progression through interplanetary space, and its interaction with the
Earth’s magnetosphere and ionosphere. The CME was an Earth-directed halo CME which
erupted following a two-ribbon M8.0 solar flare at 16:13 UT above the United States (US)
National Oceanic and Atmospheric Administration (NOAA) designated AR 10759, situated
at latitude 12° north and longitude 54°, or 11° east of the meridian, which had associated
coronal dimming later and a (possibly partial) filament eruption. Preceding this large flare
was a C1.5 class flare around 12:54 UT. There is a possibility that the CME seen in whitelight coronagraph images has two parts and thus could be made up of two separate CMEs
merging together in the low corona. The two parts appear to be travelling at different speeds
in post-processed coronagraph images. By ∼ 6 solar radii (R⊙ ), the inner brighter edge has
caught up with the outer fainter halo-like CME feature. This flare-related CME resulted in
strong coronal radio bursts (Type II and Type III) and intense MC features in the interplanetary medium near Earth. The ICME was observed as commencing at the Sun – Earth first
Lagrangian Point (L1 ) on 15 May 2005 at 02:11 UT with the Bz component of the IMF
reaching −140 nT. It impacted the Earth under an hour later and caused the most-intense
geomagnetic storm of 2005 with a Disturbed Storm Time (Dst) index reaching −263 nT at
its peak. This particular ICME was of the MC variety. Recently, it has been noted that there
were two distinct eruptions that produced two non-merging MCs (Dasso et al., 2009). In
addition, in-situ data show the presence of bi-directional electrons indicating that the field
lines remain connected to the Sun. The ionosphere responded to the storm on a global scale
as observed using a multi-instrument approach. We combine X-ray, EUV, white-light, and
interplanetary scintillation (IPS) observations as well as measurements from multiple spacecraft, ground-based instruments, and coronal/heliospheric modelling, to detail the history of
this event and summarise its features from launch to its interception by the Earth.
The first paper to discuss this event was Jones et al. (2006), who discussed the coronal radio-burst aspect of the CME. Bisi (2006) discussed some preliminary non-radial flow
analysis of the event along with some IPS velocities. Yurchyshyn et al. (2006) discusses the
initial investigations into the magnetic structure on the solar “surface” and how this compares with what is seen near Earth. Kozyreva and Kleimenova (2007) talked about geomagnetic pulsations and magnetic disturbances during the initial phase of the Earth-interaction
of the ICME/MC. Jones (2007) discussed some preliminary velocity analysis of the event
from IPS and white-light observations as well as some details of the coronal radio-burst
aspect of the CME. The outer radiation belt dynamics and some magnetospheric and nearEarth in-situ indices were discussed by Tverskaya et al. (2007). Kleimenova and Kozyreva
(2007) discussed the daytime quasi-periodic geomagnetic pulsations during the recovery
phase on to 17 May 2005. A sigmoid-shaped AR was described by Liu et al. (2007). Zhang
et al. (2007b), with their corrected CME table, provide some of the earlier CME/ICME/MCfeature values for this event. A preliminary analysis of the Carrington rotation in which this
event took place (CR2029) was carried out by Bisi et al. (2007b) using three-dimensional
(3-D) reconstruction of IPS data. Fallows et al. (2007) discussed the multi-frequency side of
IPS observations which included the 13 – 15 May 2005 CME/ICME event. This event and
The 13 May 2005 Coronal Mass Ejection
55
some preliminary discussion on its 3-D structure comparison from 3-D reconstruction to
what is seen in coronagraph data can be found in Bisi et al. (2007c). Breen et al. (2008) introduce the poleward deflection of the fast solar wind due to the passage of the 13 – 15 May
2005 CME/ICME and we also revisit this here in this article. Jain et al. (2008) noted that
there was a large decrease in cosmic-ray particles hitting the Earth around the time the event
was interacting with the Earth’s magnetosphere and this resulted in a strong Forbush Decrease (FD). An overview of ionospheric response is discussed in relation to high-Latitude
ionospheric convection and Birkeland Current response by Eriksson et al. (2008). Lara and
Borgazzi (2009) use the 13 May 2005 CME and associated Type II burst to test their analytical model for ICME transport based on hydrodynamic theory. A numerical simulation
was carried out by Xie, Gopalswamy, and St. Cyr (2009) to study the propagation of this
shock-driven event. The possibility of two non-interacting MCs resulting within the ICME
of this event is discussed by Dasso et al. (2009). Haider et al. (2009) briefly talks about
total electron content (TEC) effects at both the Earth and at Mars. Kazachenko et al. (2009)
notes that sunspot rotation nearly triples the energy for this eruption. Finally, known to date,
this event is also discussed in some detail by Gopalswamy (2010). Here, we conduct the
most-complete study to date, within a single article.
For the purpose of this article, we use the term CME to describe the 13 May 2005 event
in coronagraph white-light imagery and its initiation, ICME when seen in heliospheric observations of IPS and in our heliospheric 3-D tomographic reconstruction in interplanetary
space, and MC when describing the event’s effects and interaction with the Earth’s magnetosphere and ionosphere. We review many of the available data covering this solar eruption,
the passage of the associated ICME through the inner heliosphere and its interaction with
the space environment of the Earth. The aim throughout has been to discuss the initiation,
development, propagation, and planetary impact of the event in a clear sequential manner,
and thus to provide a secure starting point for future studies of this event. A summary of
each of the key steps in sequence of this event from Sun to Earth can be seen in here Table 1. In Section 2 we describe the observing techniques and models used in this paper in
some detail. Section 3 describes the different types of observations and measurements of the
13 May 2005 CME and its propagation from the Sun to the Earth. We form a brief overall
summary and have been able to draw some preliminary conclusions, and we discuss these
in Section 4.
This paper emphasises the need for multiple techniques to analyse these many data sets
available for studying eruptive events and their effects on the Earth and Earth environment.
These are necessary to try to understand each process, or step, along the way from the early
aspects of initiation near (or “on”) the Sun, right through to the effects on, and interaction
with, the Earth.
2. Techniques and Models
This section covers a brief description of the techniques and models we use here for the 13 –
15 May 2005 combination CME/ICME/MC event sequence, and serves as a brief overview
of each for context in Section 3.
2.1. Observations and Modelling of Active Regions, Coronal Dimming, and Coronal
Magnetic Features
Many observations of solar phenomena such as EUV observations of active regions (ARs),
coronal dimming (often as a result of flare/CME initiation), and coronal magnetic features
56
Table 1 Summary table of the various parts of the 13 – 15 May 2005 CME/ICME/MC event sequence as discussed in this paper. The magnetosphere compression was a result
of the ICME sheath crossing the Earth. Any thus-far undefined acronyms will be subsequently defined in the detailed information given in Sections 2 and 3.
13 – 15 May 2005 CME/ICME/MC Sequence and timing of each part
Description
Instrument
Value
Date and Time (UT)
Active Region (AR)
SOHO|CDS
Latitude 12°/Longitude 54°
13 May 2005
Flare
GOES 12
C1.5 Class Flare
13 May 2005 12:54
13 May 2005 16:13
Two-ribbon flare eruption
GOES 12
M8.0 Class Flare
EUV brightenings
SOHO|EIT
N/A
13 May 2005 16:37
CME initiation
Various
500 Mm height
Initiation speed
SOHO|LASCO
POS 1689 km s−1
13 May 2005 ∼ 16:42
Interplanetary transit speed
Multi-IPS
900 – 1200 km s−1
13 – 15 May 2005
Shock arrival at L1
SOHO|CELIAS
15 May 2005 02:11
ICME/Sheath arrival at L1
Wind|SWE
∼ 1100+ km s−1
Shock arrival at Earth
Geotail/Cluster 1
Magnetopause Compression/Sheath Crossing
15 May 2005 02:40
∼ 1000 km s−1
13 May 2005 17:52
15 May 2005 ∼ 02:11
PCN Max.
Magnetometer
11.1 nT
15 May 2005 03:15
Kp Max.
Global Magnetometer Network
83
15 May 2005 06:00
AU
Auroral Magnetometers
1.3 nT
15 May 2005 06:05
AL
Auroral Magnetometers
15 May 2005 06:30
Bz component max.
Geotail
−1.7 nT
Dst min.
Equatorial Magnetometers
−263 nT
15 May 2005 08:00
−140 nT
15 May 2005 06:30
M.M. Bisi et al.
The 13 May 2005 Coronal Mass Ejection
57
can be taken with instrumentation aboard the SOlar and Heliospheric Observatory (SOHO)
(Domingo, Fleck, and Poland, 1995) spacecraft in orbit about the Sun – Earth L1 . Event studies are characterised by an exhaustive investigation of carefully picked individual CMEs that
have been well observed. Event studies have required significant development of theoretical models to be able to ingest observed data and to directly simulate observed quantities.
The Predictive Science, Inc., inner-coronal model has been developed to the point where it
can now perform detailed CME studies using observed photospheric magnetic fields with
strengths approaching kilo-Gauss (kG) fields and with spatial resolutions approaching one
arc sec. With this capability, the detailed structure of the coronal magnetic field can be represented, allowing simulation of the detailed evolution of flare ribbons and to compare this
with the observed evolution. Such detailed comparisons are the best way to evaluate whether
a particular model of CME initiation is viable, and are an excellent way to study the connection between source regions on the Sun and MCs.
The following opportunities exist to compare model predictions with observations. In
the pre-event corona, the presence (or absence) of filaments, and their morphology can be
used, as can the nature of the EUV and X-ray emission, and vector-magnetic-field signatures
in the photosphere that can be measured with vector magnetograms, for example with the
SOHO – Michelson Doppler Imager (SOHO|MDI) (Scherrer et al., 1995). In the CMEinitiation and -eruption phase, it is useful to look at the evolution of Hα flare ribbons, the
morphology of hard X-ray emission kernels as observed by, for example, the Reuven Ramaty
High-Energy Solar Spectroscopic Imager (RHESSI) (see, e.g., Dennis et al., 1996), SOHO
– Extreme-Ultraviolet Imaging Telescope (SOHO|EIT) (Delaboudinière et al., 1995) waves
(EIT waves), the inferred CME shock properties as deduced from radio measurements and
energetic-particle acceleration, and in-situ magnetic-flux-rope properties.
The previous study of the 12 May 1997 CME event has demonstrated that the model
can successfully match observations of the pre-event corona, and subsequently during the
initiation phase of the CME, indicating that the model is now ready to be used on more
challenging cases (subject of a forthcoming paper).
These comparisons with observations are enhanced and refined in this study of the 13
May 2005 CME. Observations will be able to be compared in greater detail, since stronger
photospheric magnetic fields (approaching kG) can now be modelled. Whereas the 12 May
1997 CME simulations were limited in comparing with the morphology of Hα ribbons because of heavy smoothing of the magnetic field used in this first attempt; the 13 May 2005
CME models the structure of magnetic field at a greater spatial resolution. This enables the
comparison with Hα ribbons in much greater detail.
2.2. White-Light Corona and Coronal Modelling
A new technique in coronal image processing enables the separation of dynamic features
from the quiescent background corona in coronagraph observations. These are used in Section 3.1 to isolate the CME signal in SOHO – Large Angle Spectrometric COronagraph
(SOHO|LASCO) (Brueckner et al., 1995) C2 and C3 images. The technique is described
in detail by Morgan and Habbal (2009). This technique first involves the removal of the
F-corona and other stable image features (e.g. instrumental stray light) using the method of
Morgan and Habbal (2007). The sharp radial gradient in brightness is then removed using
the ‘Normalizing Radial Graded Filter’ (NRGF) (Morgan, Habbal, and Woo, 2006). Finally,
regions of the corona not containing a CME are fitted to a polynomial function of brightness,
and are removed. The final image contains isolated CME signal and noise.
NRGF images are also used as a base for solar rotation tomography (SRT), enabling
qualitative maps of the coronal density structure to be built. A series of LASCO C2 images
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M.M. Bisi et al.
covering a time period of half a solar rotation can be used to reconstruct a density map of
the corona for that period. This is different to building a standard coronagraph Carrington
or synoptic map, since tomographical techniques are used which help to resolve each line of
sight (LOS). The method used is described by Morgan, Habbal, and Lugaz (2009).
2.3. Radio Bursts and Their Propagation
Radio bursts from the solar atmosphere and solar wind display a characteristic fall in frequency with an increasing distance from the Sun. This arises from the decreasing density
(and thus plasma frequency) of the solar plasma at greater distance as the shock in the plasma
moves outward from the Sun and into the interplanetary medium. Observations of the variation of the frequency with time as the frequency decreases can provide a measure of the
velocity at which the shock is travelling outward from the Sun. The relationship between the
plasma frequency, νp , and the electron density, Ne , can be described by Equation (1) where
the plasma frequency is given in Hz:
e 2 Ne
νp =
Hz = 81Ne Hz.
(1)
2
4π ǫ0 me
Here, Ne is the electron number density in m−3 ; ǫ0 is the permittivity of free space = 8.854×
10−12 F m−1 ; e is the electronic charge of the electron = −1.602 × 10−19 C; and me is the
mass of the electron = 9.109 × 10−31 kg. Figure 1 gives an example of the radio-burst flux of
Type II and Type III radio bursts as compared with the quiet corona and that of the cosmicmicrowave background (taken from Gopalswamy, 2004).
Solar radio bursts of Type II are characterised by a narrow band of intense radiation that
drifts with time and distance from the Sun. Type II solar radio bursts are produced by the
excitation of plasma waves in the ambient medium by a shock wave propagating outward
from the Sun. These Type II emissions occur at the fundamental (F) and/or harmonic (H) of
the plasma frequency, f , which is related to the square root of the electron plasma density,
n, in the source region:
√
(2)
f = a n;
−3
here a = 9 for the F, a = 18 for the H, and f is measured in kHz and n in cm .
Figure 1 The level of
interplanetary radio-burst flux of
Type II and Type III radio bursts
as compared to the quiet coronal
emission and the cosmic
background noise (adopted from
Gopalswamy, 2004). The in-situ
plasma frequency can be seen
from the low-frequency end of
the solid curve and the
ionospheric cut-off at the
high-frequency end of the
spectrum.
The 13 May 2005 Coronal Mass Ejection
59
Type II radio bursts are classified according to their wavelength regime: metric (m), decametric/hectometric (DH), and kilometric (km) bands. It is well established that DH to km
wavelength (interplanetary) Type II radio emissions are caused by the propagation of fast
CME/shocks through the interplanetary medium (Cane, Sheeley, and Howard, 1987). The
frequency-drifting Type II radio emissions can, in principle, provide continuous tracking
of a CME/shock from the solar corona all the way to 1 AU. Therefore, we can use radio
observations to reconstruct the speed profile for a CME-driven shock.
Since the interplanetary density decreases with the square of the heliocentric distance,
the CME/shock will generate radio emissions at decreasing frequencies as it propagates
through interplanetary space. Reiner et al. (1998) showed that the radio emission from a
given source region along the shock front is expected to be organised along a straight line,
assuming that the speed of the shock is approximately constant. Hence, the slope of this line
may be determined as:
v
,
(3)
slope = √
a n0 R⊙
where v is the shock speed, n0 is the density normalised to 1 AU, R⊙ = 1.5 × 108 km, a = 9
if the emission is F, and a = 18 if the emission is H.
The observed frequency-drift rate of Type II radio emissions generated by CME-driven
shocks provides a direct measurement of the shock speed. Hence the shock speed (v) is
computed as:
√
(4)
v = (slope)aR⊙ n0 .
Type III bursts occur more frequently than Type II and are the most common flareassociated radio burst. These are narrow-band emission bursts like Type II, but are characterised by their brief duration of seconds and by their rapid change of wavelengths
from decimetres to decametres. This corresponds to a decrease in frequency from around
500 MHz to 5 MHz but such Type III bursts have been observed at frequencies as low as
0.1 MHz to over 1000 MHz. Type III radio bursts can be attributed to beams of electrons
thrown out from the Sun with kinetic energies of 10 keV to 100 keV and speeds of up to
around half that of the speed of light. Type II and Type III bursts can occur immediately
one after the other, or even simultaneously. It is generally accepted that the radio waves are
emitted near the electron plasma frequency or its harmonics.
2.4. Interplanetary Scintillation (IPS) and Heliospheric Modelling
Observations of interplanetary scintillation (IPS) have been used to study the solar wind
for more than 40 years (see, e.g., Hewish, Scott, and Wills, 1964; Coles and Rickett, 1976;
Kojima and Kakinuma, 1987; Rickett, 1992; Bisi et al., 2008, 2010b). IPS arises from the
scattering of radio waves by small-scale (∼ 150 km) density inhomogeneities in the solar
wind flowing approximately radially outward from the Sun (Ne ). It is a powerful tool to
probe the interplanetary medium. If the variations in apparent source intensity of a distant,
compact astronomical radio source (e.g. a quasar) as seen by two telescopes are recorded
simultaneously at a time when the ray paths from the radio source to the antennas lie in
a plane which passes through the centre of the Sun, then the time lag for maximum cross
correlation of the signals can be used to provide a first estimate of the solar wind outflow
speed (see, e.g., Armstrong and Coles, 1972; Bourgois et al., 1985; Breen et al., 2006; Bisi
et al., 2006, 2010c, in this issue). This is only possible when the points of closest approach
of the lines of sight from the radio source to each antenna lie close to the same Sun – Earth –
source plane (see, e.g., Breen et al., 1996b; Coles, 1996). The point of closest approach of the
60
M.M. Bisi et al.
LOS to the Sun is referred to as the P-Point (or sometimes as the “impact parameter”). The
accuracy to which the solar wind speed can be determined improves as the separation of the
antennas projected into the plane of the sky (POS) increases (see, e.g., Klinglesmith, 1997;
Bisi, 2006). More sophisticated methods fitting the observed auto- and cross-correlation
spectra with the results from a weak-scattering model have also been adopted for IPS data
analyses (see, e.g., Coles, 1996; Klinglesmith, 1997; Bisi et al., 2007a; Fallows, Breen, and
Dorrian, 2008).
Observations of IPS made using the radio arrays of the Solar-Terrestrial Environment
Laboratory (STELab/STEL), Nagoya University, Japan at 327 MHz (Kojima and Kakinuma, 1987) have been routinely carried out for over 20 years. The mainland antennas
of the European Incoherent SCATter radar (EISCAT) (Rishbeth and Williams, 1985) have
been used for such observations since 1982 (see, e.g., Bourgois et al., 1985) at frequencies
around 935 MHz, and on a regular basis since 1990 (see, e.g., Breen et al., 2002) at frequencies based around 928 MHz to 930 MHz, and 1420 MHz following upgrade work on
two of the three sites (Wannberg et al., 2002). More recently, the EISCAT Svalbard Radar
(ESR) (Wannberg et al., 1997) has been used for multi-frequency IPS analyses along with
the mainland EISCAT telescopes (see, e.g., Fallows et al., 2006) operating at 500 MHz.
Extremely-long baseline (ELB) IPS observations (Bisi et al., 2005, 2007b; Bisi, 2006;
Breen et al., 2006) have also been undertaken between northern Scandinavia and the United
Kingdom (UK) through the use of the Multi-Element Radio-Linked Interferometer Network
(MERLIN) (Thomasson, 1986) radio telescopes based around the UK which have been operated at frequencies of around 1420 MHz, 5 GHz, and 6 GHz (see, e.g., Jones, 2007) for
observations of IPS (only 1420 MHz is used for the ELB IPS analyses however). The Ootacamund (Ooty) Radio Telescope (ORT) (see, e.g., Swarup et al., 1971) has also been used for
many years for observations of IPS at 327 MHz (see, e.g., Manoharan and Ananthakrishnan,
1990; Manoharan et al., 2000) as well as the large Pushchino Radio Astronomy Observatory
at ∼ 111 MHz near Moscow, Russia (see, e.g., Shishov et al., 2008).
ELB IPS presently uses both the EISCAT and MERLIN radio telescopes, which allow
for much larger separation in observing antennas (up to ∼ 2000 km parallel baselines). This
results in an increased ability to resolve multiple streams crossing the IPS ray path (LOS)
from the antenna to the radio source on the sky. Using this technique combined with a
three-mode weak-scattering IPS solar-wind model (Bisi et al., 2007a; Fallows, Breen, and
Dorrian, 2008) an extension of the originally-developed two-mode model (Coles, 1996),
three separate velocities can be fitted and obtained for material crossing the IPS LOS. In
addition, the greater-length baseline between the antennas also gives exceptional solar-wind
directional information for each of the three velocity structures crossing the ELB IPS ray
paths and these too can be fitted with the three-mode model (see, e.g., Breen et al., 2008) or
with a geometrical technique (see, e.g., Bisi et al., 2005; Bisi, 2006).
As is done in this paper with STELab IPS data, the inner heliosphere can be reconstructed in three dimensions from observations of IPS using a time-dependent 3-D tomographic model which combines many lines of sight to undo the LOS integration (see, e.g.,
Jackson et al., 2003; Jackson and Hick, 2005). For 3-D density reconstructions however,
the IPS scintillation level requires a “conversion” to g-level which subsequently provides
a proxy for density. 3-D velocity reconstructions can be carried out from the derived IPS
radial velocity values alone. Presently, the University of California, San Diego (UCSD) 3-D
reconstructions incorporate a kinematic model. The resulting reconstructions are of an innerheliosphere region typically ranging from 15 R⊙ (the “source surface”) out to approximately
3 AU. Scintillation-level measurements have been obtained using the STELab radio antenna
at Kiso from 1997, and more recently (from mid-2002) from the Fuji antenna (see Jackson
The 13 May 2005 Coronal Mass Ejection
61
and Hick, 2005; Bisi et al., 2009a; and references therein). Another STELab antenna is located at Sugadaira. The newly-operating Toyokawa site will extend the STELab system to
observe many more radio sources per day and is expected to be fully operational some time
in 2010. STELab typically observes 20 to 40+ astronomical radio sources per day at present
using the original IPS arrays, but with the old Toyokawa antenna decommissioned.
The disturbance factor (normalised scintillation level), or g-level, g, is defined by Equation (5):
g = m/m.
(5)
Here, m is the observed scintillation level, and m is the mean level of scintillation for
the source at the elongation at the time of observation. The scintillation-level measurements
from STELab are available for each astronomical radio source as an intensity variation of the
signal strength. Since density values along the LOS are not known a priori, the g-level proxy
for density is related to small-scale density variations by Equation (6) where the small-scale
density variations are assumed to have a power law dependence on heliospheric density and
heliocentric distance:
Ne = Ac R α Neβ .
(6)
Here, Ac is a proportionality constant, R is the radial distance from the Sun, α is a power
of the radial falloff, and β is the power of the density. Ac , α, and β are determined using
best-fit comparisons with in-situ data at 1 AU. The values used here are as previously used
for 327 MHz IPS observations: Ac is set equal to 1, and the two powers α and β, are −3.5
and 0.7 respectively. Further discussion of determining g-level and density can be found in
Jackson et al. (1998, 2003), and references therein.
The 3-D reconstructions from the STELab data used here have latitude and longitude
digital resolutions of 20° × 20°. The one-day time cadence has 6-hour interpolated increments, to yield output four times a day for the modelled structure of the inner heliosphere.
The accuracy of this model is very well confirmed both in this low-resolution form, and in
higher-resolution forms such as for IPS data from Ooty (see, e.g., Bisi et al., 2007c, 2009b)
or white-light observations from the Solar Mass Ejection Imager (SMEI) (see, e.g., Jackson
et al., 2006, 2008; Bisi et al., 2008). The output of these tomographic reconstructions have
successfully provided a “source surface” input into the ENLIL 3-D magneto-hydrodynamic
(MHD) numerical model (see, e.g., Odstrcil and Pizzo, 2002) which when propagated out
through the interplanetary medium, also compares reasonably with in-situ measurements
(see, e.g., Bisi et al., 2008).
Observations of IPS at 327 MHz using Ooty can probe the solar wind in the heliocentricdistance range of ≈ 10 R⊙ to ≈ 250 R⊙ (see, e.g., Manoharan et al., 2000). In the case of
the CME event on 13 May 2005, the monitoring of IPS has been made on a grid of about
500 – 600 radio sources. The calibration of the temporal power spectrum of scintillation
on each source provides the velocity of the solar wind and scintillation index along the
LOS to the source (see, e.g., Manoharan and Ananthakrishnan, 1990). Furthermore, the
scintillation index of a radio source measured at a given distance from the Sun is normalised
using its expected long-term average to obtain the normalised scintillation index (g), which,
as mentioned previously, is directly related to the level of density turbulence in the solar
wind (Manoharan, 1993). For example, the g-level value close to unity corresponds to the
ambient solar wind. Whereas g > 1 and g < 1, correspond to enhanced and depleted density
turbulence conditions of the solar wind respectively. This is the same as with other IPS glevel values. Therefore, a map of g-level produced in an interval of time when an ICME is
crossing the IPS field of view (FOV), would result in an excessive turbulence being produced
62
M.M. Bisi et al.
by the ICME and its shock would likely provide the location of the ICME with respect to
the Sun’s centre. This is shown for Ooty data in Section 3.2.
When the separation of the receiving stations for a two-site IPS observation (e.g. with
EISCAT) is sufficiently large (typically > 220 km), separate fast and slow peaks become resolved in the cross-correlation function (CCF) for scintillation detected at the two telescopes
(Breen et al., 1996b). Coles (1995, 1996) showed that these functions could be best interpreted by treating the LOS ray path for the observation as containing discrete regions of fast
and slow wind, each with its own set of characteristic properties, an interpretation confirmed
by the results from the Ulysses (Fossi, Noci, and Poletto, 1992) first orbit (see, e.g., Riley
et al., 1997; McComas et al., 1998b). By mapping the ray path ballistically back down to
the corona and overlaying it on a map of white-light intensities derived from coronagraph
observations over a solar rotation, it was possible to associate faster flow with those sections
of the ray path overlying dark coronal regions and slow with those above bright regions,
and thus use a weak-scattering model to fit for the fast and slow components separately
(see, e.g., Breen et al., 1996a). More recently, this approach has been extended to produce a
three-mode weak-scattering model (as previously mentioned), allowing stream-interaction
regions (SIRs) and solar wind transients to be investigated in a more realistic manner.
In analysing the EISCAT – MERLIN ELB observation of the 13 May 2005 CME, a
slightly different approach was adopted in constraining the fit of the three-mode model:
instead of projecting the ray path back to a “coronal source surface” and attempting to estimate the angular extent of the CME about its presumed source region (as was done by, for
example, Jones et al., 2007), we took advantage of the 3-D reconstructions of solar wind normalised scintillation levels derived from STELab IPS measurements. The LOS ray path for
the ELB IPS observation was overlaid on the reconstructed density in 3-D space, allowing
the regions of the LOS immersed in unperturbed fast wind, fast wind near the CME boundary, and the CME itself, to be identified, and these three components of flow fitted separately
(Breen et al., 2008). This reduction in uncertainty in the locations of these flow components
greatly improved our confidence in the accuracy of the resulting velocity estimates from the
three-mode model.
2.5. Near-Earth In-situ (Interplanetary) Measurements
CMEs remove plasma and magnetic field from the Sun expelling them out into interplanetary space where they later become ICMEs. Measurements made by a single in-situ spacecraft provide some information about the local magnetic topology, and measurements of
suprathermal and more energetic particles indicate the Parker spiral or a loop within the
ejecta, plasma composition, and connectivity of the ICME fields back to the Sun. Even so,
this information is obtained by measurements along a radial vector to the Sun as the ICME
material passes over the spacecraft.
MCs, as previously mentioned in Section 1, are generally considered a subset of ICMEs
(Gosling, 1990) are characterised by an enhanced magnetic-field strength, a smooth and
large rotation of the magnetic-field vector, and low proton temperature. The 3-D structures of
MCs have been successfully determined using flux-rope inversion models (Goldstein, 1983;
Marubashi, 1986; Burlaga, Lepping, and Jones, 1990). Various in-situ studies with multiple
spacecraft have shown that these magnetic-rope models are, however, insufficient to determine the true cross-sectional dimensions, orientation, and non-force-free nature of typical
flux-rope ICMEs (Russell and Mulligan, 2002). The in-situ magnetic-field measurements
are carried out with spacecraft instrumentation such as the Advanced Composition Explorer
(ACE) (Stone et al., 1998) – MAGnetometer (ACE|MAG) (Smith et al., 1998).
The 13 May 2005 Coronal Mass Ejection
63
CMEs are frequently associated with filament eruptions. The directions of the MC axes
are found to be roughly aligned with the disappearing filaments (Bothmer and Schwenn,
1994, 1998). However, it is not so easy to quantify this association, and further developments
are needed to really understand the fundamental transport mechanisms and interaction with
the ambient solar wind (see, e.g., the review by Démoulin, 2008). Thus, views of these
structures on their way from the Sun to 1 AU are extremely important in the determination
of the overall configuration of material outflow. The connection between the solar “surface”
magnetic configuration of the 13 May 2005 features of this event to those at 1 AU have
already been attempted by Dasso et al. (2009).
2.6. Magnetosphere and Ionosphere Measurements
Highly-detailed measurements of electromagnetic fields and particle distributions can be
obtained from in-situ instrumentation aboard spacecraft. These measurements are local to
each particular spacecraft. When studying the impacts of solar ejecta on the Earth’s magnetosphere (and indeed any process that affects the magnetosphere on a global scale), however,
a synoptic, multi-instrument approach combining space- and ground-based measurements
from various sources is necessary to put those detailed measurements into context.
Here we use data from the Geotail (Nishida, 1994), Cluster (Escoubet, Schmidt, and
Goldstein, 1997; Escoubet, Fehringer, and Goldstein, 2001), Los Alamos National Laboratory (LANL), and the Imager for Magnetopause-to-Aurora Global Exploration (IMAGE)
(Gibson et al., 2000) spacecraft along with ground-based data from the NORSTAR Meridian
Scanning Photometer (MSP) array (Rostoker et al., 1995) and geomagnetic indices derived
from various ground-based magnetometers to describe the impact of the 13 May 2005 CME
on the Earth’s magnetosphere.
A typical measure of a geomagnetic storm’s effect on the Earth is given by the Equatorial
Dst; this is derived from the horizontal component of the magnetic variations observed at
low latitudes. This is caused by the equatorial ring current which flows westward at all local
times during the main phase of large geomagnetic storms.
Studies of particular relevance to the present work are those of Ebihara et al. (2005b,
2005a), and Hanuise et al. (2006). In the former study, the storm was observed and the Dst
index, based upon the magnetic disturbance associated with the high-latitude convection
pattern, reached a maximum negative excursion of −472 nT. This indicated that the antisunward cross-polar flow was significantly enhanced. The high-latitude convection pattern
expanded equatorward to 60° MLAT. The response of the high-latitude ionosphere to an
ICME was shown by Hanuise et al. (2006). The Auroral Electrojet (AE) index, a measure
of the variation in the horizontal component of the geomagnetic field in the auroral zone
increased, suggesting the presence of AEs and the polar cap potential, a measure of the electric field driving the high-latitude convection pattern also rose, which indicated enhanced
convection across the polar cap.
The Polar Cap Index (PCI) is an index for magnetic activity based upon the magnetic
disturbance associated with the high-latitude convection pattern (Troshichev et al., 1988).
This disturbance, when measured close to the pole, indicates the anti-sunward flow at all
local times. Measurements of the PCI in the northern hemisphere (PCN) are made at Thule,
Greenland.
The AE Indices AU and AL are derived from the variations in the horizontal component
of the geomagnetic field observed in the auroral zone in the northern hemisphere. They indicate the magnitude of the AE. The AU and AL indices are intended to express the strongest
current intensity of the eastward and westward AEs respectively. For each station a base
64
M.M. Bisi et al.
value is calculated for each month from the average of all data from the five quietest days
on a global scale. This value is subtracted from each observation. AU and AL are the largest
and smallest values observed at a given time after processing. The geomagnetic index, Kp,
is derived from a global network of observatories and provides an indicator of the global
geomagnetic disturbance.
3. The Observations and Measurements of the 13 – 15 May 2005 CME/ICME/MC
This section covers our observations, interpretations, and modelling of the 13 – 15 May 2005
CME/ICME/MC. It is split into three sub-sections and subsequent parts within them, in a
similar context to that of Section 2. The 13 May 2005 CME was selected as the subject of an
event study by a National Aeronautics and Space Administration (NASA) Living with a Star
(LWS) Focused Science Team on the “CME-ICME Connection”. It had good observations,
an associated magnetic flux rope, and it also had a more complex photospheric and coronal
magnetic field than the earlier-mentioned 12 May 1997 CME event. In addition, it is a US
National Science Foundation (NSF) sponsored Solar, Heliospheric, and INterplanetary Environment (SHINE) community event which had three working group sessions during the
SHINE-based Meetings in 2008, 2009 and in 2010.
We first provide a brief overview of Carrington rotation 2029 (CR2029) in which this
CME occurred. There were a vast number of events during this time, which is consistent
with the period being a few years into the declining phase from solar maximum. From the
Association of Lunar & Planetary Observers (A.L.P.O.) Solar Section Carrington rotation
commencement dates listings (http://alpo-astronomy.org/solar/rotn_nos.html), the CR2029
start date/time was 21 April 2005 07:52:27.84 UT, and the CR2029 end date/time was 18
May 2005 13:29:08.15 UT. From the Coordinated Data Analysis Workshops (CDAW) CME
catalog http://cdaw.gsfc.nasa.gov/CME_list/, there were some 116 events noted as being seen
in the SOHO|LASCO white-light coronagraph data. Thus, much of interplanetary space was
likely dominated by transient activity throughout most (if not all) of CR2029.
Reconstructed/modelled synoptic-map overviews of CR2029 can be seen in Figures 2
and 3 using three different methods. The first (Figure 2) shows a reconstructed white-light
Carrington map of SOHO|LASCO data at 4 R⊙ using the technique briefly described in
Section 2.2 and in detail by Morgan, Habbal, and Lugaz (2009). The density is in normalised (arbitrary) units (see Morgan, Habbal, and Lugaz, 2009). The second (Figure 3) is
a comparison of 3-D reconstructed IPS results using the UCSD kinematic, time-dependent
solar-wind model with STELab IPS data (top) along with a 3-D, time-dependent coronal
MHD model (bottom). The MHD Model of the Solar Corona is also known as the MAS
model. The top plot is produced from our IPS results at a height from the Sun of 30 R⊙
and is an average of each of the individual daily reconstructions of solar wind structure in
the inner heliosphere using STELab IPS data as input. The bottom plot is a velocity profile
produced using the results from the MHD model, driven by the observed LOS photospheric
magnetic field. The details of the algorithm have been discussed by Mikić et al. (1999) and
Linker et al. (1999), and references therein. Its extension from the solar corona to the inner heliosphere is discussed by Riley, Linker, and Mikić (2001) and Riley et al. (2001). We
computed the speed at 30 R⊙ using the coronal magnetic configuration: at 1 R⊙ we set the
radial speed to be some value, vslow , at the boundary between open and closed field lines
over a width of ∼ 6° (in a direction normal to the boundary) and smoothly raised it to vfast
over ∼ 3°. We then mapped this speed profile outward along the open field lines to 30 R⊙ .
Although this may appear somewhat ad hoc, it is based on the commonly-held view that fast
The 13 May 2005 Coronal Mass Ejection
65
Figure 2 Reconstructed Carrington map at 4 R⊙ from post-processed SOHO|LASCO imagery. This map
shows the large-scale structure of the corona at this height. Sunspots are shown as red crosses superimposed
on the image. The NOAA AR 10759 (Latitude N12°/Longitude 54°) lies just poleward (north) of what appears
to be an edge of the streamer. The density scale gets increasingly dense from black through purple to blue,
then green, and finally through yellow and orange to red for the most-dense portions in arbitrary units.
wind emanates from within coronal holes and slow wind is associated with the boundary between open and closed fields; as would be the case if closed field lines were sporadically
opened through magnetic reconnection. Comparisons with in-situ measurements suggest
that this approach is capable of reproducing the essential features of the large-scale structure of the inner heliosphere for a variety of solar conditions (Riley, Linker, and Mikić, 2001;
Riley et al., 2001). There are several areas where the comparison does reasonably well, but,
as can be expected, there are also some significant disagreements. The STELab IPS results
show speeds approaching 1000 km s−1 at times. This is unlikely to come from the ambient
solar wind, based on Ulysses measurements at high latitudes (see, e.g., Riley et al., 1997)
using the Solar Wind Observations Over the Poles of the Sun (Ulysses|SWOOPS) instrumentation (Bame et al., 1992). These high values are likely associated with the multiple
CMEs that occurred throughout CR2029. However, the structure of higher-than-Ulysses velocities within the fast solar wind (away from transients) has been thoroughly studied by
Bisi et al. (2007a, 2010a) using EISCAT IPS data. The two low-speed blobs at longitudes
of 240° and 330° match up reasonably well. The general sense of the low-speed band tracks
well in longitude, but the IPS band is considerably thinner than the MHD solution. At the
higher latitudes, the IPS results show considerably more structure.
Figure 4 shows SOHO|EIT images taken close to the initiation time of the event and
then about an hour after initiation displaying a long-lived brightening of the AR and coronal
dimming to the south and toward Sun centre. This in itself hints that there may have been
more than one initiating event. Figure 5 shows two images from SOHO|LASCO C2 and C3
coronagraphs of the event. Further details can be found in the respective captions.
66
M.M. Bisi et al.
Figure 3 Reconstructed Carrington maps at 30 R⊙ from STELab IPS observations (top) and the MAS MHD
propagation model (bottom). A more complete description can be found in the text.
Figure 4 13 May 2005 SOHO|EIT images at 16:57 UT (left) and 17:37 UT (right), taken from Bisi et al.
(2007b). The NOAA AR 10759 responsible for this flare/CME (bright area circled) along with associated
dimming region (dark area circled). This was a relatively long-lasting period of activity for this AR. Images
originally taken from the CDAW CME Catalog.
The 13 May 2005 Coronal Mass Ejection
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Figure 5 SOHO|LASCO images of the 13 May 2005 CME, adapted from Bisi et al. (2007b). The top pair
of images uses the standard LASCO processing techniques and are originally taken from the CDAW CME
Catalog. The bottom pair of images use the ‘Normalizing Radial Graded Filter’ (NRGF) post-processing
as described in Section 2.2. The Halo CME pictured at 17:22 UT in LASCO C2 (left) and at 17:42 UT in
LASCO C3 (right) with an estimated LASCO plane-of-the-sky (POS) speed of 1689 km s−1 (as stated in the
CDAW CME Catalog). The CME first C2 appearance was on 13 May 2005 at 17:12:05 UT and the CME first
onset at 1 R⊙ was at 16:47:34 UT. Notice the double loop-like structure (circled) to the east of the Sun – Earth
line in both sets of images and a faint near-halo feature (at all position angles in the post-processed images)
seen more prominently in the C3 images. This image comparison displays the enhancement to the features
due to the NRGF processing as compared with the standard processing. Further study with these improved
images can be found in Section 3.1.
3.1. The Sun and Corona
As mentioned, a strong brightening can be seen in SOHO|EIT 195 Å at 16:37 UT on 13 May
2005, with many smaller precursor events throughout this day. The result was a fast CME,
with an estimated POS speed of 1689 km s−1 leaving the Sun. It had a pre-eruption sigmoid in the NOAA Geostationary Operational Environmental Satellites (GOES) 12 Solar
X-Ray Imager (SXI) (Hill et al., 2005; Pizzo et al., 2005) and dimming regions in EIT.
68
Table 2 Table showing a
summary of the key initiating
event fluxes; further details are
found in the text.
M.M. Bisi et al.
Key Event Initiation Fluxes
Description
Value
Injected magnetic flux
4.4 × 1022 Mx
Magnetic-field strength overlying filament
6G
Drag coefficient
1.15
Flux rope ellipse eccentricity
0.87
Prominence density
2.1 × 108 cm−3
RHESSI hard X-ray data, Transition Region And Coronal Explorer (TRACE) (Wolfson
et al., 1997) data, radio, and Solar Energetic Paricle (SEP) data are also available. The source
AR was well observed with both SOHO|MDI, and the Big Bear Solar Observatory (BBSO)
(http://www.bbso.njit.edu/) vector magnetogram. The data from these are all presently available, thus providing an overall good magnetic coverage of the solar surface. We note, though,
that the quality of the vector magnetogram is questionable for direct quantitative use.
The linear force-free model for the coronal field above AR 10759 (Abramenko and
Yurchishin, 1996) has a left-handed helicity with the axis (poloidal) field oriented toward
heliographic north. The eruption-process properties resemble a tether-cutting model (Moore
and Labonte, 1980; Moore et al., 2001) according to Yurchyshyn et al. (2006) and Liu et al.
(2007). Using measurements from SOHO|LASCO, ACE|MAG, and ACE – Solar Wind
Electron, Proton and Alpha Monitor (ACE|SWEPAM) (McComas et al., 1998a) instrument
data as boundary conditions, Yurchyshyn et al. (2006) fit the parameters of the Krall et al.
(2006) erupting-flux model determining the injected magnetic flux necessary to produce the
magnitude of the magnetic field at 1 AU, the overlying magnetic-field strength stabilising the
filament, the drag coefficient of the solar wind, the eccentricity of the flux rope ellipse, and
the initial density of the cold plasma in the prominence. These are summarised in Table 2.
To study the solar surface energetics, a background corona and solar wind with an empirical heating was first created in preparation for studies of CME initiation. Figure 6 shows
a comparison between the simulated and observed EUV emission. It is a reasonable first
attempt, though there are significant differences that will be improved over time. These additional refinements will likely be the subject of future publications.
We started exploring the energisation of the AR magnetic field for this event and decided
to energise the AR solely by emerging transverse field from below the photosphere. No
shear was applied when modelling this event, since we noticed from the 12 May 1997 study
that transverse magnetic-field emergence was very effective in producing fast eruptions. The
emergence was accomplished in an ad hoc manner by specifying the appropriate bounding
electric fields using masks and functions related to traces of the photospheric neutral line,
and the distance away from it. The details were presented by Z. Mikić at the 2009 SHINE
Meeting working group on this event in Nova Scotia, Canada (Mikić et al., 2009).
These preliminary simulations were done at high spatial resolution with kG magnetic
fields using a zero-beta version of the model. This flux emergence produced highly-sheared
field at the neutral line and forms a filament channel. The morphology of the filament channel is very similar to the observed Hα filament using the Improved Solar Observing Optical
Network (ISOON) (http://nsosp.nso.edu/isoon/), as shown in Figure 7. This indicates that
the ad hoc technique is reasonable. In Figure 8, the radial and transverse magnetic fields in
the photosphere after the emergence are shown, and these are compared with the starting potential field. It can be seen that there is considerable shear at the neutral line. This modelled
The 13 May 2005 Coronal Mass Ejection
69
Figure 6 A quantitative comparison between observed (top) and simulated (bottom) emission for 13 May
2005. This coronal state was used for our CME-initiation studies. The parameters in the heating model may require some fine-tuning to improve the match, especially within ARs where the simulated emission is presently
overly bright at all heights.
field can be compared directly with the magnetic field measured with a vector magnetogram
to check how reasonable the energisation mechanism is.
The SOHO – Coronal Diagnostic Spectrometer (SOHO|CDS) (Harrison et al., 1995)
instrument observed and followed the NOAA AR 10759 from 13 to 15 May 2005. Observations were recorded during two time intervals on 13 May 2005: from 07:40 UT to 15:08 UT
for the first set; and later from 17:56 UT to 23:29 UT for the second set. The data employed were taken using the FLARE_AR (ID 173 var. 9) observing sequence. This takes
five selected wavelength ranges in the EUV portion of electromagnetic spectrum from both
the Normal Incidence Spectrometer (NIS) channels. For NIS1: 359.56 Å – 361.95 Å; and
for NIS2: 582.34 Å – 586.33 Å, 590.19 Å – 594.18 Å, 622.92 Å – 626.92 Å, and 627.27 Å –
631.15 Å. The sequence rasters cover an area of the Sun of 183 arc sec by 181 arc sec. Each
raster takes approximately 11 minutes to complete. The first set of observations includes
39 rasters, while the second set provides an additional 30 rasters. Table 3 lists some of the
observation details.
The five emission lines observed are shown in Table 4. The neutral helium line (He I)
at 584.33 Å, is a strong line formed in the upper chromosphere/lower transition region.
It views the coolest plasma and is in-effect a chromospheric indicator. The oxygen line
(O V) at 629.73 Å, is a bright line at transition region temperatures, and the magnesium
line (Mg X) at 624.94 Å, is a quite-bright coronal line. The iron line (Fe XVI) at 360.76 Å,
is used to determine high-temperature plasmas and is particularly bright in ARs but not
usually bright in the quiet Sun. Finally, the other iron line (Fe XIX) at 592.23 Å, identified
as a magnetic-dipole transition within the ground configuration (Widing, 1978), is a typical
signature of flares. The formation temperature of each line was taken from the peaks of the
contribution functions calculated using the Atomic Data and Analysis Structure (ADAS)
database (Summers, 2004; http://adas.phys.strath.ac.uk/). Thus in summary, the five lines
70
M.M. Bisi et al.
Figure 7 Comparison between an ad hoc generated filament and an Hα Improved Solar Observing Optical
Network (ISOON) image for the 13 May 2005 CME event study. The simulation was energised by emergence
of transverse magnetic field from below the photosphere in the vicinity of the neutral line. The ISOON
image, which was taken about two minutes after the commencement of the flare, shows the beginning of the
flare ribbons. The filament has already erupted partially, but there is good agreement between the simulated
filament and the ISOON filament, and the filament channel delineated by the flare ribbons. Images (a) and
(b) show the Hα image with model field lines superimposed, while (c) shows the detail of the Hα image
of the AR. The panels in (d) show the development (from left to right) of the sheared field in the filament
channel as the transverse field has emerged.
cover a large temperature range simultaneously (3 × 104 K to 9 × 106 K), allowing us in
principle, to examine the contribution of the ejection processes of plasmas at very different
temperatures in the solar atmosphere.
The SOHO|CDS data for the two sets of observations were reduced and calibrated (Lang
et al., 2002). Only three individual times from each set were chosen, and plots of intensity
The 13 May 2005 Coronal Mass Ejection
71
Figure 8 Simulated “vector magnetograms” showing the magnetic field in the photosphere. The radial component of the magnetic field is shown in the intensity image; red indicates positive values and blue indicates
negative values. The arrows show the transverse magnetic field. The neutral line is indicated in black. The
left panel (a) shows the potential magnetic field that is used to initiate the calculation. The right panel (b)
shows the magnetic field after transverse field has emerged in the vicinity of the neutral line. The emerged
field creates a filament channel with highly-sheared magnetic field along the neutral line.
and spectra were computed. These can be seen in Figures 9 through 12. Figures 9 and 10
give the distribution of intensity of the examined AR before the eruption, which occurred
at 16:13 UT (Figure 9), and in the post-flare phase (Figure 10). The sigmoid shape (see,
e.g., Canfield, Hudson, and McKenzie, 1999) that characterises this AR before the flare and
subsequent CME (as described in some detail by Liu et al., 2007) could be seen more than
eight hours before the event (images at 07:40 UT in Figure 9). The symmetric shape is well
defined in the Fe XVI and Fe XIX bands (such as in the EIT images at 171 Å and 195 Å, and
in the TRACE 171 Å images), but is also clearly visible in the Mg X and He I bands, where
the emission is more intense near the footpoints. After the explosive M8.0 flare event, CDS
observed the AR from 17:56 UT until 23:29 UT showing the evolution from the sigmoidal
field to loops of an arcade (as illustrated by Figure 10). The spectra, presented in Figures 11
and 12, describe the change in the averaged intensity for each wavelength band, both in
the pre-flare and post-flare state (respectively) providing a general overview of the event. In
addition, the integrated intensity along the line profiles is plotted as function of time, taking
into account the full sets of rasters (Figure 13). The increase of intensity just before the
flare is not evident because of the lack of CDS observations when the eruption happened
(between 15:08 UT and 17:56 UT). However, it is possible to detect a signature of such
increase in the plot for the second set of observations, where, after a maximum value, the
intensity of each line starts to decrease rapidly. An earlier detection of the smaller flare may
also be possible from these spectra and plots.
These basic spectroscopic tools and qualitative information may allow for an analysis
of the CME onset process and the subsequent phase by detection of plasma flows in the
various temperature regimes, with supporting magnetograms and observations from coronagraphs.
72
M.M. Bisi et al.
Table 3 Details of the observations gathered using SOHO|CDS on 13 May 2005.
First Set
Filename
Observation
Exp.
X-cen.
Y-cen.
X-width
Y-width
Object
Time (UT)
s
(arc sec)
(arc sec)
(arc sec)
(arc sec)
(NOAA)
s32622r00
07:40:30
8.00
s32622r02
08:01:47
8.00
−262.4
s32622r03
08:13:19
8.00
s32622r04
08:24:51
8.00
s32622r05
08:36:25
8.00
s32622r06
08:47:55
8.00
s32622r07
08:59:27
8.00
s32622r08
09:10:58
8.00
s32622r09
09:22:30
8.00
s32622r10
09:34:05
8.00
s32622r11
09:45:34
8.00
s32622r12
09:57:07
8.00
s32622r13
10:08:38
8.00
s32622r14
10:20:10
8.00
s32622r15
10:31:45
8.00
s32622r16
10:43:14
8.00
s32622r17
10:54:46
8.00
s32622r18
11:06:18
8.00
s32622r19
11:17:51
8.00
s32622r20
11:29:24
8.00
s32622r21
11:40:54
8.00
s32622r22
11:52:26
8.00
s32622r23
12:03:58
8.00
s32622r24
12:15:30
8.00
s32622r25
12:27:04
8.00
s32622r26
12:38:34
8.00
s32622r27
12:50:06
8.00
s32622r28
13:01:38
8.00
s32622r29
13:13:10
8.00
s32622r30
13:24:45
8.00
s32622r31
13:36:14
8.00
s32622r32
13:47:46
8.00
s32622r33
13:59:18
8.00
s32622r34
14:10:50
8.00
s32622r35
14:22:24
8.00
s32622r36
14:33:54
8.00
s32622r37
14:45:26
8.00
s32622r38
14:56:58
8.00
s32622r39
15:08:30
8.00
234.2
182.8
181.4
AR 10759
−262.2
234.1
182.8
181.4
AR 10759
−262.2
234.0
182.8
181.4
AR 10759
−262.4
234.2
182.8
181.4
AR 10759
−250.1
232.9
182.8
181.4
AR 10759
−250.5
233.3
182.8
181.4
AR 10759
−250.6
233.4
182.8
181.4
AR 10759
−250.5
233.5
182.8
181.4
AR 10759
−250.6
233.3
182.8
181.4
AR 10759
−245.6
238.2
182.8
181.4
AR 10759
−245.5
238.1
182.8
181.4
AR 10759
−245.7
238.3
182.8
181.4
AR 10759
−245.9
238.5
182.8
181.4
AR 10759
−245.8
238.4
182.8
181.4
AR 10759
−239.3
231.9
182.8
181.4
AR 10759
−239.3
231.9
182.8
181.4
AR 10759
−239.3
231.9
182.8
181.4
AR 10759
−239.3
231.9
182.8
181.4
AR 10759
−239.3
231.9
182.8
181.4
AR 10759
−226.9
234.7
182.8
181.4
AR 10759
−226.8
234.6
182.8
181.4
AR 10759
−226.8
234.6
182.8
181.4
AR 10759
−226.9
−226.9
−217.5
−217.6
−217.8
−217.5
−217.6
−209.0
−208.9
−208.9
−208.9
−208.9
−204.2
−204.1
−204.2
−204.2
−204.1
234.7
182.8
181.4
AR 10759
234.7
182.8
181.4
AR 10759
234.7
182.8
181.4
AR 10759
234.5
182.8
181.4
AR 10759
234.4
182.8
181.4
AR 10759
234.7
182.8
181.4
AR 10759
234.5
182.8
181.4
AR 10759
235.5
182.8
181.4
AR 10759
235.6
182.8
181.4
AR 10759
235.6
182.8
181.4
AR 10759
235.6
182.8
181.4
AR 10759
235.6
182.8
181.4
AR 10759
230.7
182.8
181.4
AR 10759
230.6
182.8
181.4
AR 10759
230.8
182.8
181.4
AR 10759
230.7
182.8
181.4
AR 10759
230.6
182.8
181.4
AR 10759
The 13 May 2005 Coronal Mass Ejection
73
Table 3 (Continued)
Second Set
Filename
Observation
Exp.
X-cen.
Y-cen.
X-width
Y-width
Object
Time (UT)
s
(arc sec)
(arc sec)
(arc sec)
(arc sec)
(NOAA)
s32624r00
17:56:47
8.00
s32624r01
18:07:12
8.00
−179.5
s32624r02
18:18:04
8.00
s32624r03
18:29:36
8.00
s32624r04
18:41:08
8.00
s32624r05
18:52:42
8.00
s32624r06
19:04:12
8.00
s32624r07
19:15:45
8.00
s32624r08
19:27:16
8.00
s32624r09
19:38:48
8.00
s32624r10
19:50:23
8.00
s32624r11
20:01:52
8.00
s32624r12
20:13:24
8.00
s32624r13
20:24:56
8.00
s32624r14
20:36:28
8.00
s32624r15
20:48:02
8.00
s32624r16
20:59:32
8.00
s32624r17
21:11:04
8.00
s32624r18
21:22:36
8.00
s32624r19
21:34:08
8.00
s32624r20
21:45:42
8.00
s32624r21
21:57:12
8.00
s32624r22
22:08:44
8.00
s32624r24
22:31:48
8.00
s32624r23
22:20:16
8.00
s32624r25
22:43:22
8.00
s32624r26
22:54:52
8.00
s32624r27
23:06:24
8.00
s32624r28
23:17:56
8.00
s32624r29
23:29:28
8.00
234.5
182.8
181.4
AR 10759
−179.5
234.5
182.8
181.4
AR 10759
−179.5
234.5
182.8
181.4
AR 10759
−179.5
234.5
182.8
181.4
AR 10759
−179.5
234.5
182.8
181.4
AR 10759
−168.2
232.5
182.8
181.4
AR 10759
−168.2
232.5
182.8
181.4
AR 10759
−168.1
232.6
182.8
181.4
AR 10759
−168.1
232.6
182.8
181.4
AR 10759
−168.1
232.6
182.8
181.4
AR 10759
−160.0
234.9
182.8
181.4
AR 10759
−160.1
234.7
182.8
181.4
AR 10759
−160.2
234.8
182.8
181.4
AR 10759
−160.3
234.7
182.8
181.4
AR 10759
−160.2
234.8
182.8
181.4
AR 10759
−154.2
229.4
182.8
181.4
AR 10759
−154.1
229.5
182.8
181.4
AR 10759
−154.2
229.4
182.8
181.4
AR 10759
−154.2
229.4
182.8
181.4
AR 10759
−154.1
229.5
182.8
181.4
AR 10759
−142.6
233.6
182.8
181.4
AR 10759
−142.5
233.5
182.8
181.4
AR 10759
−142.7
−143.2
−142.8
−135.8
−135.8
−135.8
−135.8
−135.8
233.7
182.8
181.4
AR 10759
234.2
182.8
181.4
AR 10759
233.8
182.8
181.4
AR 10759
234.4
182.8
181.4
AR 10759
234.4
182.8
181.4
AR 10759
234.4
182.8
181.4
AR 10759
234.4
182.8
181.4
AR 10759
234.4
182.8
181.4
AR 10759
The front edge of the 13 May 2005 CME event is observed by the SOHO|LASCO C2
coronagraph at 17:22 UT, then by the C3 coronagraph at 17:42 UT. Both these observations
are shown using appropriate post-processing (as described in Section 2.2) in Figure 14.
Unfortunately, due to a special observing program running during this day, and due to an
emergency stop in observations as SOHO was being bombarded by SEPs, these are the only
decent images available of this CME. This seriously limits the amount of information that
can be gained about the CME close-in to the Sun.
As mentioned previously, the source of the CME is AR 10759, situated at latitude
12° north and longitude 54° (or 11° east of the meridian). If the CME propagates along
a close-to-radial path from the AR, this can explain why the brightest parts of the CME as
74
Table 4 Summary table of the
SOHO|CDS emission lines
observed during 13 May 2005.
M.M. Bisi et al.
Transition Region Lines
Ion
λ (Å)
log(Te (K))
Atomic Transition
He I
584.33
4.45
OV
629.73
5.30
1s 2 1 S0 − 1s2p 1 P1
1s 2 2s 2 1 S0 − 1s 2 2s2p 1 P1
Coronal Lines
Ion
λ (Å)
log(Te (K))
Atomic Transition
Mg X
624.94
6.05
Fe XVI
360.76
6.35
1s 2 2s 2 S1/2 − 1s 2 2p 2 P1/2
Fe XIX
592.23
6.95
2p6 3s 2 S1/2 − 2p6 3p 2 P1/2
2s 2 2p4 3 P2 − 2s 2 2p4 2 D2
seen in the LASCO image are inclined somewhat to the north-east corona. Despite being a
halo CME, it is not centred directly along the Sun – Earth direction. Most clearly seen in the
C2 image of Figure 14 (left), there are two main parts to this halo CME. The first part is a
leading edge which completely surrounds the Sun, resembling a skirt with nearly constant
brightness at all angles, with the exception of a faint region to the south. The second part
is a far brighter edge limited to the north-east sector. In the C2 image, the bright edge is
distinct from the leading fainter front. In the C3 image, the bright edge has caught up with
the outer edge. Other structural details of interest are the large voids leading and following
the north-east bright edge, and a very complicated, tangled, small-scale structure within the
main body of the CME.
It is difficult to determine a speed for such a complicated structure as this CME. The
fact that only two images show the CME with any clarity does not help. Different parts of
the CME seem to travel with different speeds, and the analysis is limited by the biggest
problem in coronal observations, the LOS. The true 3-D extent of the CME in the inner
corona cannot be determined from these two images taken from the same direction, therefore
the true velocity of the CME cannot be determined accurately. By measuring the apparent
speed of various structures within the CME, we can provide a safe lower limit to the CME
velocity. Figure 15 shows the profile of the processed LASCO brightness at four different
position angles marked on the images of Figure 14. The profiles taken along position angles
35° and 75° in the C2 image are dominated by the bright inner edge close to a height of
4 R⊙ . For both these profiles, there is an enhancement on the leading (largest height) part
of this bright peak. This is the position of the fainter outer skirt of the CME. The other
two profiles at 220° and 280° show the fainter outer skirt of the CME at a similar height of
4.5 R⊙ . In the LASCO C3 plot of Figure 15 (bottom plot), the bright edge in the north-east
has overtaken the fainter skirt at position angle 35° and a height of 6.5 R⊙ . At 75°, the bright
edge has nearly caught up with the outer skirt near a height of 5.7 R⊙ . From comparing the
profiles for C2 and C3, we calculate an apparent (or POS) speed of around 840 km s−1 for
the fainter outer skirt, and around 1000 km s−1 for the inner brighter front. Since we know
that the CMEs are not propagating in the POS, the true speed is considerably higher. These
values are also significantly lower than the CDAW CME Catalog POS speed of 1689 km s−1
(as also used later in this paper for velocity-profile studies).
Are the two distinct structures a single CME structure or two separate CMEs which happen to coincide in the LASCO images due to their different speeds? Lugaz et al. (2009)
talk of CMEs which propagate at different speeds cannabilising each other in interplanetary
space, and maybe the same is happening in this case – the faster, brighter CME in the northeast catches up with the slower, fainter halo CME and interacts. Jones et al. (2007) also dis-
The 13 May 2005 Coronal Mass Ejection
75
Figure 9 A sequence of 182.8 – 181.4 arc sec images of the NOAA AR 10759 taken by SOHO|CDS on
13 May 2005 before the M8.0 eruptive flare event which occurred at 16:13 UT the same day. Here, three
different times have been selected: 07:40 UT, 11:29 UT, and 15:08 UT. These pre-flare EUV images exhibit a
very symmetric S-shape (called sigmoid) structure related to a magnetic configuration with a high probability
of eruption to produce flares and possible associated CMEs (Canfield, Hudson, and McKenzie, 1999). Such
a shape is well defined in the Fe XVI and Fe XIX band images, but also clearly visible in the Mg X and He I
bands; even if the emission is more intense near the footpoints, which can be better localised by the help of
the O V images.
76
M.M. Bisi et al.
Figure 10 As in Figure 9, but at different times. Again, another three times have been selected: 17:56 UT,
20:48 UT, and 23:29 UT. These SOHO|CDS image sequences, taken following the explosive M8.0 flare
event, show the evolution from the sigmoidal field (Figure 9) to loops of an arcade.
cuss this possibility of CME/ICME cannibilisation through detailed studies of CMEs/ICMEs
using EISCAT observations of IPS. In support of this latter scenario is the fact that the faint
outer CME seen in Figure 14 is a pure halo CME, only slightly skewed to the north-east
(consistent with a CME propagating along latitude 12° north and longitude 10° east, radial
from AR 10759). The second brighter CME appears only in the north-east, and is different
The 13 May 2005 Coronal Mass Ejection
77
Figure 11 A sequence of averaged spectra derived from the first set of SOHO|CDS observations (see Figure 9 caption). The plots use the same y-range scaling for each line both in the first (here) and in the second
set of observations (Figure 12). It more clearly shows how the line intensities change during the 13 May 2005
events. A very small increase in the 11:29 UT flux may suggest very faint and early signs of the C1.5 flare
which occurred at 12:54 UT.
in appearance to a halo CME. The bright front edge extends down towards the Sun at both
ends, which is inconsistent with a halo CME. AR 10759 is very dynamic throughout much of
13 May 2005; EIT was observing every 20 minutes or so in the 195 Å bandpass. Figures 16
and 17 show a series of images of the AR and surrounding area. The main “culprit” for the
78
M.M. Bisi et al.
Figure 12 A sequence of averaged spectra derived from the second set of SOHO|CDS observations (see
Figures 10 and 11). The big increase in the average intensity at 17:56 UT compared to the set shown in
Figure 11 especially in the highly-ionised Fe XIX band, which is the signature that a huge explosive event
occurred (in this case, the M8.0 flare). The post-flare phase is indicated by the intensity decay for all the
wavelength bands at the subsequent times.
CME event is the activity after 16:00 UT shown in Figure 17. A bright flare dominates the
AR at this time. If this event, at around 16:15 UT, results in the CME seen in LASCO C2
at 17:22 UT, the average apparent (or POS, or minimum) speed of the CME will be around
500 – 600 km s−1 . However, earlier in the day at 11:25 UT, the same AR ejects material.
The 13 May 2005 Coronal Mass Ejection
79
Figure 13 Intensity – time curves for the emission lines observed by SOHO|CDS (first set of observations:
upper plot; second set of observations: lower plot). The intensity is in photons s−1 cm−2 sr−1 . The time is
in hours starting from 00:00 UT on 13 May 2005. The upper plot of the earlier observations seems to show a
small increase of emission for the majority of the lines, which coincides with the timing of the C1.5 flare and
could be as a result of that flare around 12:54 UT.
There is no evidence of any CME activity in the LASCO data until the 17:22 UT observation. Therefore, the first eruption from the AR at 11:25 UT does not result in a CME, or the
CME is narrow and is heading directly towards Earth (so is not seen past the occulting disks
of LASCO). It is highly unlikely for this first eruption from the AR to have resulted in the
CME activity viewed by LASCO at 17:22 UT. This does not negate the possibility that the
earlier brightening is not related to the C1.5 flare however.
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Figure 14 (Left) Post-processed LASCO C2 image of 13 May 2005 at 17:22 UT and (right) LASCO C3
image taken at 17:42 UT on the same day. The C2 field of view (FOV) shown here extends from 2.5 R⊙ to
5.95 R⊙ , whilst the C3 FOV extends from 5 R⊙ to 10 R⊙ . The inner black circles toward the centre of each
image show the position of the Sun’s limb. Processing details can be found in the text. The dashed red lines at
position angles 35°, 75°, 220°, and 280° are the lines along which profiles are plotted later in Figure 15 and
described in the text.
Figure 15 Profiles of normalised brightness along the dashed red lines shown in Figure 14. The top plot is
for LASCO C2 at time 17:22 UT, the bottom for C3 at 17:42 UT (both on 13 May 2005). From these plots it
is possible to estimate an approximate POS speed for the CME(s) as described in the text.
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Figure 16 Four consecutive zoomed-in images of the region close to AR 10759 as observed by the EIT
telescope in the 195 Å bandpass. The four observations are made at times (from left to right) of 11:05 UT,
11:25 UT, 11:42 UT, and 12:05 UT. AR 10759 sits in the top left of these images. The second image shows a
tongue of material ejected out from the AR towards the bottom of the image. In the third image, there is still
ejecta material apparent. The third image also shows signs of dimming. Again, signs of the earlier C1.5 flare
perhaps.
Figure 17 As with Figure 16 but for times (from left to right across the top followed by left to right across
the bottom) of 16:09 UT, 16:26 UT, 16:37 UT, 16:57 UT, 17:07 UT, 17:27 UT, 17:37 UT, and 22:57 UT.
A bright flare dominates the AR in the third image until the seventh. The last image at 22:57 UT shows the
result of the dramatic reconfiguration of the coronal magnetic structure during this day.
3.1.1. A Brief Overview of the Sun and Corona Observations and Measurements
The post-processed LASCO coronagraph images show a halo-like structure coming toward
the Earth, but with a bias to the north-east quadrant. This holds some strength to the possibility that there was a secondary eruption, which quickly merged with the first in the low
corona. This is also a possible agreement with the two MCs detailed in Dasso et al. (2009).
With regards to the possible earlier C1.5 flare having involvement with the CME/ICME
initiation, this is indeed possible (Dasso et al., 2009 also touch on this). From the
SOHO|CDS data, three peaks could be seen in the time series of intensities (Figure 13),
especially in the oxygen-line and helium-line intensities. The first two, before 12:00 UT,
with the biggest around 13:00 UT, may be peaks which tell us something about the flare
– the upper plot of the earlier set of observations is a good signature of the C1.5 flare that
happens at 12:54 UT. Of course, these data show clearly the latter parts of the much larger
M8.0 two-ribbon flare, and this is seen in the lower time series plot of intensities in this same
figure.
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Figure 18 Observations of radio signals from EISCAT and MERLIN showing the anomalous signal-strength
enhancement on 13 May 2005 due to the flare/CME and associated coronal radio burst overwhelming the IPS
signal as taken from Jones et al. (2006). All EISCAT and MERLIN observations took place at 1420 MHz
with the exception of one (as marked here in the figure).
3.2. The Inner Heliosphere
Here, we discuss the 13 – 15 May CME/ICME/MC event as the ICME and associated
shock/sheath were detected and progressed through the inner heliosphere from outside of
coronagraph fields of view through radio-burst tracking and observations of IPS, to near its
measurement in situ around the Sun – Earth L1 point.
3.2.1. Coronal Radio Bursts
The coronal radio burst associated with the 13 May 2005 flare/CME was first detected by the
EISCAT and MERLIN antennas while conducting observations of IPS (Jones et al., 2006)
as shown in summary in Figure 18. An off-axis response created an enhancement in power
received at each of the antennas. Recordings were made simultaneously with the GOES12 satellite and the PHOENIX-2 radio spectrometer (Messmer, Benz, and Monstein, 1999)
as seen in Figure 19. Also marked on Figure 18 are the start, peak, and end times of the
GOES-detected flare showing very good correlation with the detections made by EISCAT
and MERLIN. These are assumed to be the initial detections of the Type III radio burst as
a result of the flare/CME eruption. Precise details on how this result was concluded can be
found in Jones et al. (2006).
The direct radio observations discussed here were acquired by the Wind – Radio and
Plasma Waves Experiment (Wind|WAVES) (Ogilvie and Desch, 1997; Bougeret et al.,
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Figure 19 Figure displaying the PHOENIX-2 radio spectrogram recorded on 13 May 2005 from 16:30 UT
onwards as taken from Jones et al. (2006). The clear onset of enhanced emission at lower frequencies from
16:35 UT can be seen in the image, which later led to a strong radio burst from 16:45 UT. The evolution from
high-frequency to lower-frequency emission is a characteristic of a coronal radio burst as the resultant blast
wave from the explosive event, in this case a flare/CME, moves outward from the Sun through the corona.
This is explored further in the text.
1995). The Wind|WAVES experiment has several sensitive radio receivers: Radio Receiver Band 1 (RAD1), Radio Receiver Band 2 (RAD2), and the Thermal Noise Receiver
(TNR); these cover a frequency range from 4 kHz to ∼ 14 MHz. One-minute resolution
data from the RAD1 receiver were analysed and a new analysis technique, using the frequency drift of DH-km-Type-II (DH-km-TII) spectra obtained by the RAD1, applied to
infer, at some adequate intervals, the propagation speed of the IP shock associated with
the ICME. The technique is explained in detail elsewhere (Aguilar-Rodriguez et al., 2005;
Gonzalez-Esparza and Aguilar-Rodriguez, 2009).
On 13 May 2005 at 17:12 UT, SOHO|LASCO observed a fast (POS v = 1689 km s−1 )
halo CME as previously described. This CME was associated with the M8.0 soft X-ray flare.
The flare region was localised close to the centre of the disk, at N12 E11. Near the Sun, a
metric Type II burst was seen in radio emission detected by the Wind|WAVES experiment
(Figure 20a) and reported as fundamental/harmonic emission by the Radio spectrometer of
Astrophysikalisches Institut Potsdam, starting on 13 May 2005 at ∼ 16:38 UT and ending
the same day at ∼ 16:50 UT, covering a frequency range of ∼ 350 MHz to 40 MHz (see
Figure 20b). The event was observed at ∼ 17:00 UT on 13 May 2005 at 5 MHz, drifting
slowly down to 40 kHz at 02:10 UT on 15 May as reported in the Wind|WAVES online list
(http://lep694.gsfc.nasa.gov/waves/waves), maintained by M.L. Kaiser.
We began the analysis by isolating the Type II event from the RAD1 dynamic spectrum
by setting to zero any emission outside the Type II burst feature. Figure 21a shows this
RAD1 spectrum along with 21b showing the Type II burst isolated from the dynamic spectrum with a vertical line indicating a time cut (23:29 UT). Then each RAD1 spectrum was
analysed to obtain its spectral properties. Figure 21c shows the flux density versus the fre-
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Figure 20 (a) Metric Type II radio burst observed by the Radio spectrometer of Astrophysikalisches Institut Potsdam on 13 May 2005. (b) Interplanetary counterpart of the Type II radio burst detected by the
Wind|WAVES experiment using the RAD1+RAD2 receivers.
quency of the Type II radio burst at 23:29 UT (marked using asterisks’) and a Gaussian fit to
the flux-density distribution (dashed line). This Gaussian approximation to the flux-density
distribution gives us, when there was a good fit, the central frequency (dashed – dotted line),
which is the fundamental emission line in the radio spectrum at the given time. This procedure cleans the data from emission contamination. Once the central frequency emissions
were obtained, the frequency drift associated with the DH-km-TII radio emission was studied. Figure 21d shows the central emission frequency versus time for the CME on 13 May
2005. The time interval is about 21 hours. Although the frequency shows a tendency to de-
The 13 May 2005 Coronal Mass Ejection
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Figure 21 Central frequency drift and shock speed deduction. (a) WAVES/RAD1 dynamic spectrum associated with the propagation of the CME/shock. (b) Isolated Type II solar burst (dotted line) temporal cut
to perform the central frequency analysis. (c) Flux density versus frequency for 1-min RAD1 spectrum at
23:29 UT. The Gaussian fit gives the central frequency. (d) Central frequency versus time. The horizontal
lines show the four intervals to perform the Type II slope analysis. (e) Shock speed results applying the
Type II slope analysis in all the sub-intervals within interval 2. The shock speed and its uncertainty are given
by the mean value and the standard deviation of the distribution of speed values.
crease with time, the drift is not smooth and there are regions where the frequency increases.
We selected four intervals with a well-defined drift to infer the slope of the frequency and
using Equation (3) (previously described in Section 2.3), choosing a value for n0 and if it
is the F or H emission, we obtained the shock speed. The frequency slope was obtained
by applying a linear fit to the central frequency points in the interval. The slope was then
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Table 5 Summary table of the primary observations of IPS for the 13 – 15 May 2005 CME/ICME/MC event.
Individual observations of IPS for the 13 – 15 May 2005 CME/ICME/MC
IPS Observing System
Source Observed
Frequency of Observation
Dates Observed
13 – 14 May 2005
EISCAT
J0319+415
928 MHz/1420 MHz
MERLIN
J0319+415
1420 MHz
13 – 14 May 2005
Ooty
J0319+415
327 MHz
14 May 2005
Ooty
J0352-071
327 MHz
14 May 2005
Ooty
J0158-141
327 MHz
15 May 2005
Ooty
J0447-220
327 MHz
15 May 2005
Pushchino
J0137+331
111 MHz
13 – 16 May 2005
Pushchino
J0043+520
111 MHz
14 – 16 May 2005
STELab
J0323+053
327 MHz
13 – 14 May 2005
STELab
J0336+322
327 MHz
13 – 14 May 2005
STELab
J0431+204
327 MHz
14 May 2005
STELab
J0534+193
327 MHz
13 – 14 May 2005
calculated for every set of 60 (90 – 120) consecutive central frequency points in the interval.
This gives an ensemble of speed solutions associated with all the sub-intervals. Figure 21e
shows the speed solutions for interval 2. The distribution of points shows that, in this case,
the speed solutions tend to converge towards the maximum value. We define the DH-km-TII
shock speed (1230 km s−1 ) as the mean of the distribution, and its associated uncertainty
(30 km s−1 ) as the standard deviation of the distribution.
3.2.2. Observations of IPS
A summary of the primary individual IPS radio sources observed relating to the 13 – 15 May
2005 CME/ICME/MC can be seen in Table 5. Also shown are the IPS-observing systems
that each radio source was observed along with the observing frequency, date of the observations, and also the source name in J2000 coordinates. One of the most-well-known IPS
sources, as observed here by EISCAT, MERLIN, and Ooty, is that of J0319+415 (3C 84).
The individual STELab observations of IPS seen in Figure 22 show velocity sky maps
across the top panels and g-level sky maps across the bottom. These are summaries on
the sky with the Sun at the centre and angles measured outward in all directions (fisheye
projections); the first marker out from the centre corresponds to 0.5 AU (= sin 30°), and the
second corresponds to 1 AU (= sin 90°). No completely-clear signature is seen in the g-level
values for the 13 – 15 May 2005 ICME, but there are clear enhancements in the velocity
values. Several observations showed velocities over 800 km s−1 with marked increases in
the same (or similarly-located IPS sources on the sky) observed from the 13 May 2005
through 14, some of which continued again on the 15 and 16 May. Although some g-level
values are greater than unity (and hence suggesting increased turbulence and/or relatively
increased density), none can be directly associated with the ICME. This is where the 3-D
reconstruction technique can improve on this, as discussed later in this section and described
in Section 2.4.
Time series of the Ooty velocity observations around this period can be seen in Figure 23.
Similarly, the Ooty g-level observations of IPS can be seen in Figure 24 as plots on the sky
displaying the regions where enhanced scintillation can be seen (otherwise known as g-map
displays).
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Figure 22 The STELab velocity (top) and g-level (bottom) results of IPS observations displayed as sky
maps, as if the Sun is at the centre of each for 13 through 16 May 2005 (from left to right). There are several
high-velocity values during this time, particularly around 14 – 15 May covering nearly all of the sunward sky,
plus a large g-level increase to the East on 14 May. Although this is not conclusive evidence alone, it does
suggest that a higher-velocity flow (which is possibly more highly-turbulent material) is coming toward the
Earth from the Sun and may be biased in magnitude East of the Sun – Earth line. Further explanations are
found in the text with regards to how these plots fit into the other supporting IPS observations of this ICME.
Figure 23 Ooty velocity estimates are plotted as a function of time for observations taken on 13 and 14
May 2005. The upper and lower plots are of the same data. However, in the upper plot only a 15-minute
average of the P-Point distance of the observations is displayed. The west and east points shown by red and
blue colour-coded triangles indicate observations that have been taken to the western and eastern sides of the
Sun – Earth line respectively. As can be seen by this figure and also by Figure 24, the SOHO|LASCO halo
CME’s ICME is slightly oriented to the east side of the Sun. This is consistent with the STELab results and
also the 3-D reconstructed STELab IPS data displayed later in Figure 29. In addition, this agrees with the
earlier-discussed SOHO|LASCO post-processing analysis in Section 3.1.
At Ooty, on 14, 15, and 16 May 2005, four scans of the sky were observed per day, i.e. two
scans to the west and two to the east of the Sun, all in a distance range of ≈ 10 R⊙ to 250 R⊙ .
Figure 23 displays velocity estimates obtained from Ooty on 13 – 14 May 2005. Each value
on the lower plot indicates the IPS observation corresponding to a radio source. The P-Point
distance of the LOS is plotted in the upper plot. As shown in the figure, the observations from
Ooty in a day from the morning to late night (i.e. in the UT time range 20:00 UT on 13 May
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Figure 24 This figure includes g-maps from Ooty observations of IPS on 14, 15, and 16 May 2005 for
the indicated time periods. In these “position angle – heliocentric distance” plots, the Sun is at the centre.
The concentric circles are of distance of 50 R⊙ , 100 R⊙ , 150 R⊙ , and 200 R⊙ . As can be seen, the halo
ICME is expanding in the east – west direction. In each plot, the time increases from west to east (refer to
Figure 23). For example, in the top-left plot, the halo ICME onset can be seen at the east of Sun, in the heliocentric-distance range of 100 R⊙ – 125 R⊙ from north to south latitude range. In the next time-range observations (top-right plot), the ICME has moved outward as seen in the western side of the Sun at about 175 R⊙ .
As a greater number of sources are observed from west to east of the Sun, the ICME can be seen in east of the
Sun more strongly. On the same day, at later time (bottom-left plot). Now, the ICME has moved to > 200 R⊙
distance range. In the last plot (bottom-right plot), the ICME goes beyond Earth’s orbit.
2005 to 18:00 UT on 14 May 2005) scan the interplanetary medium twice, from west of the
Sun to the east. Before the ICME onset in the IPS window, ∼ 20:00 UT on 13 May 2005, the
speed range was between 300 km s−1 and 450 km s−1 ; thus indicating the background solar
wind speed present. When the ICME crossed the IPS FOV (IPS sources in the transition
region from the west of the Sun to the east), the velocity for sources passing through the
ICME increased to between 900 km s−1 and 1100 km s−1 during the time interval 01:00 UT
to 04:00 UT on 14 May and in the distance range of 80 R⊙ to 90 R⊙ . The next west – east
scan (07:00 UT to 18:00 UT time range), which included different sources, also showed
nearly similar conditions for the ICME crossing. The nose of the ICME is oriented to the east
of the Sun and it is consistent with the CME onset seen in the SOHO|LASCO FOV. On the
next day, ∼ 06:00 UT on 15 May 2005 (see Figure 24), the ICME moved to ∼ 150 R⊙ and
in this distance range, the velocity dropped to between around 800 km s−1 and 1000 km s−1 .
On 16 May 2005, the ICME had crossed the IPS FOV and the velocity estimates went back
to the background solar-wind values of around 300 km s−1 to 500 km s−1 .
Figure 24 shows the g-maps of different IPS fields of view between 14 and 16 May 2005.
In these “PA – Distance” images on 14 May 2005, the onset of the ICME can be seen in the
heliocentric-distance range of about 100 R⊙ to 125 R⊙ . The halo shape of the ICME can be
clearly seen in the west – east direction of the Sun around 150 R⊙ on the image for the time
interval of between 19:00 UT and 06:30 UT on 14 – 15 May 2005. In the next consecutive
image, 10:00 UT to 18:00 UT on 15 May 2005, the ICME has propagated to a heliocentric
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Figure 25 Positions of the
J0137+331 (3C 48) and
J0043+520 (3C 20) radio sources
in ecliptic coordinates as
observed by the Pushchino Radio
Astronomy Observatory.
distance of ∼ 200 R⊙ and is more dominant to the east. In the last scan, the ICME has
already moved out of the Earth’s orbit and beyond the Ooty IPS FOV.
The Pushchino Radio Astronomy Observatory was used at the frequency of 111 MHz
and a sampling rate of 0.1 s with a frequency bandwidth of 600 kHz for IPS observations of
several distant radio sources during May 2005. The most-interesting data obtained were for
two of these observed sources: J0137+331 (3C 48, B0134+329, CTA 15) and J0043+520
(3C 20, B0040+517, CTB 6). The ecliptic coordinates of the sources are equal to λ48 = 36°,
β48 = 23°, λ20 = 34°, and β20 = 43° for J0137+331 and J0043+520, respectively. The
longitude of the Sun on 14 May 2005 was λ0 = 53°, correspondingly λ20 − λ0 = 17°, λ48 −
λ0 = 19°. The locations of the sources relative to the Sun are shown in Figure 25. One can
see that two observations were sampling the solar wind plasma in close proximity. However,
due to the strong difference in angular sizes of the sources, the IPS signals are sensitive to
plasma turbulence at different regions of the solar wind. The radio source J0137+331 having
angular size about 0.1 arc sec at this frequency, can be considered as a “point source”. The
region of effective radio-wave intensity modulation on the LOS to this source is defined by
the inequality in quiet conditions:
0.6 AU = 1 AU(cos ε48 − 0.6 sin ε48 ) < z < 1 AU(cos ε48 + 0.6 sin ε48 ) = 1.2 AU.
(7)
Here, z is the distance from the observer to the irregularities, ε48 ≈ 25° is the source elongation angle on 14 May 2005, the LOS P-Point was located at the heliocentric distance of
around 0.4 AU, and lower and upper slab boundaries correspond to the 1/2 turbulence level
in comparison with that of the P-Point. Radio source J0043+520 is extended; its angular size
is several arc sec. Large angular size results in suppression of the contribution from the distant region and approaching the effective region to the observer. The effectively modulating
region of the turbulent plasma is bounded in quiet conditions by
0 < z < 1 AU cos ε20 ≈ 0.7 AU.
(8)
Here, ε20 ≈ 45° is the elongation angle on 14 May, and the LOS P-Point is located at heliocentric distance of about 0.7 AU. The geometry of the IPS experiment at 111 MHz in the
heliocentric coordinate system can be seen in Figure 26.
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M.M. Bisi et al.
Figure 26 Geometry of lines of
sight from the observations of
IPS of radio sources J0137+331
(3C 48) and J0043+520 (3C 20)
by the Pushchino Radio
Astronomy Observatory.
Effectively modulating regions
are marked by the solid thicker
portions of each LOS.
Figure 27 Daily scintillation
indices for the two IPS sources
J0137+331 (3C 48) and
J0043+520 (3C 20) observed by
the Pushchino Radio Astronomy
Observatory.
Calculated values of scintillation indices for source J0137+331 (3C 48) – root mean
square (rms) source flux fluctuations normalised to the mean flux – are shown in Figure 27
in logarithmic scale. A weak increase of about 1.5 times in scintillation index was observed
on 14 May for source J0137+331. Corresponding temporal scintillation spectra are presented on the left-hand side of Figure 28. It follows from the comparison of the three spectra
that no significant variations of the spectrum width, and correspondingly, of the moving
scintillation pattern speed were observed. The source J0137+331 observed at 12:20 UT on
14 May was nearly at closest approach to the Sun, its LOS P-Point was projected to the
north solar hemisphere at latitude of about 50° and at nearly 60° to the east from central
meridian. Weak IPS increase for this source shows that a disturbance passed through the
LOS at the very periphery of the effective slab, marked in Figure 26, or even closer to the
Earth. This is similar to the EISCAT – MERLIN observations of J0319+415. Absence of a
distinct increase in the plasma velocity is explained by the projection effect, because only
transversal to the LOS component of the velocity produces the IPS pattern motion. Using
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91
Figure 28 Temporal IPS power spectra of the source J0137+331 (3C 48 – left) and J0043+520 (3C 20 –
right) as observed using the Pushchino Radio Astronomy Observatory.
the time delay between the event in the corona and the J0137+331 IPS observation, one can
obtain only the lower estimate of the mean ICME speed ≈ 800 km s−1 .
In quiet conditions for source J0043+520 (3C 20), the IPS response is usually weakly
scintillating with a scintillation index typically around 0.03. Figure 27 shows that strong,
about three times that, enhancements of IPS caused by the 13 May 2005 CME were observed
for the source J0043+520 on 14 and 15 May 2005. This source is sensitive to large-scale
disturbances in the sector of heliographic longitudes of 0° to 45° to the east from central
meridian. The value of scintillation index on 16 May again corresponds back to a quiet IPS
signal. Temporal IPS power spectra of J0043+520 are presented in the right-hand panels of
Figure 28.
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M.M. Bisi et al.
The weak IPS increase for J0137+331 and considerable increase for J0043+520 in the
geometry of the presented observations (Figure 26) allows us to draw the conclusion that the
ICME disturbance in the eastern part of the northern hemisphere was strong only within the
angle about 50°. The IPS spectrum on 16 May 2005 is typical for the big source with angular
size about 3 – 5 arc sec in quiet solar-wind conditions (Shishov and Shishova, 1979). A combination of radial dependence of turbulence power and the source-size truncating effect,
results in the same scintillation spectra as scintillations on a statistically quasi-uniform extended slab adjacent to the observer (Shishov and Shishova, 1979). Temporal IPS spectrum
on 14 May is qualitatively similar to that on 16 May, but has about 10 times increased integral power and has shifted to higher scintillation frequencies. Frequency shift corresponds
to a three times increase in the local plasma speed. The temporal scintillation power spectrum measured on 15 May (as in Figure 28) is similar to the spectrum of a point source and
differs strongly from the spectra on 14 and 16 May, especially in the steep high-frequency
region for source J0043+520. Such a spectrum of an extended source, as well as scintillation
index increase and typical scintillation time decrease, are explained by a narrow slab with
enhanced turbulence located within the distance 0.2 AU from the observer. Local plasma
velocity was approximately the same as in undisturbed conditions. The approaching of the
effective slab to the Earth results in reduction of the source-size influence on the scintillation
parameters and their similarity to the parameters of a point source. IPS data for the source
J0043+520, observed on 14 May at 13:20 UT, show that the ICME front passed the distance
more or about 0.7 AU for 21 hours with the mean speed ≈ 1300 km s−1 . The same value of
the local plasma speed is needed for explanation of three times typical IPS time decrease by
the passing ICME. Near the Earth, a slab with enhanced turbulence on 15 May at 13:20 UT
is associated evidently with the tail part of ICME because the front part has already reached
the Earth at around 03:00 UT resulting in the large geomagnetic storm. Time delay of about
13.5 hours between J0043+520 on 14 May 2005 and the event in the Earth magnetosphere
corresponds to the mean ICME speed ≈ 900 km s−1 between 0.7 AU and 1 AU heliocentric
distances, showing that the ICME decelerated inside 1 AU. These are in agreement with the
other IPS results and those discussed in Section 3.2.1 on the radio-burst propagation, and
more strongly show that the ICME was perhaps more biased in the north-east quadrant.
Next we look at the EISCAT – MERLIN cross-correlated observations of IPS source
J0319+415 (3C 84) in the context of the 3-D tomographic reconstruction of the inner heliosphere from the STELab IPS data, before moving on to looking at the reconstruction itself
in greater detail. Figure 29 shows the 3-D reconstructed ICME on 14 May 2005 at 03:00 UT
using the methods described in Section 2.4 and taken from Breen et al. (2008). The fast solar
wind to the north of the ICME received a strong perturbation non-radially northward in the
meridional direction due to the interplanetary transit of the 13 – 15 May 2005 ICME (Bisi,
2006; Breen et al., 2008). Non-radial flows of, and possibly due to CMEs/ICMEs, have also
been noted by Owens and Cargill (2004) through other methods. In Figure 29, the 3-D reconstructed density structure from the STELab observations (during CR2028-CR2030) on
14 May 2005 at 03:00 UT can be seen along with the EISCAT – MERLIN ELB IPS LOS
ray path passing through the reconstructed volume. The LOS is immersed in three different structures in the heliosphere from very near the Earth (equatorial intermediate/fast solar
wind) through towards the P-Point (the ICME itself) and passed the P-Point out to around
2 AU from Earth (the polar fast solar wind stream). It is the latter of the three that dominate
this observation of IPS with the ICME only producing a small contribution to the received
scintillation pattern at the EISCAT and MERLIN receivers on 14 May 2005 (as described in
greater detail in Breen et al., 2008). The fast polar wind stream is also nearest to, and intersecting the largest portion around (particularly on the source end), the LOS P-Point, which
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Figure 29 3-D reconstruction of the distribution of reconstructed solar-wind density upwards of 8 e− cm−3
(brighter colours toward yellow mean increasing density) at 15:00 UT on 14 May 2005 as derived from
STElab IPS data. All non-ICME-related features have been removed for clarity of viewing when displaying
the 3-D volume. The axes are heliographic coordinates with X-axis direction pointing toward the vernal
equinox, and Z-axis directed toward solar heliographic north. An r −2 density fall-off has been removed
from the analysis to better-show structures further out from the Sun, which in this image, is represented by
the central yellow/orange sphere and the Earth by the blue sphere (not to scale). The EISCAT – MERLIN
IPS LOS ray path observation of J0319+415 (3C 84) on 14 May 2005 is shown as a line going from the
Earth towards the upper-left section on the image. This is cutting through the reconstructed Earth-directed
13 – 15 May 2005 ICME. The density structures intersecting and crossing the IPS LOS, which the three-mode
weak-scattering model fit suggests, are associated with the main, second, and third peaks in the observed
CCFs (as described in the text), and are marked as “fast wind diverted polewards (main peak)”, second
peak is unlabelled but is caused by the solar wind at the Earth end of the LOS nearest the ecliptic, and
“CME (small peak)” respectively. Further details on the EISCAT – MERLIN fit can be found in the text along
with the relevant references. Adapted from Breen et al. (2008), reproduced by permission of the American
Astronomical Society (AAS).
is where the greatest amount of scintillation power is created (see, e.g., Fallows, 2001). It is
also this stream which is highly perturbed meridionally northward by ∼ 8° – 15° because of
the passage of the ICME.
The EISCAT – MERLIN velocities were obtained from the observations of J0319+415
on 13 and 14 May 2005. This source is a relatively-high-latitude source to the east of the
Sun – Earth line (where the CME/ICME appears to dominate both in the SOHO|LASCO
images, in the 3-D reconstructed results from the STELab IPS data, and in the other observations of IPS already described).
The EISCAT – MERLIN observations on 13 May 2005 showed a clear fast-wind velocity
up until 16:20 UT – 16:30 UT, when strong radio interference (interpreted by Jones et al.,
2006 as a coronal noise burst, detected in the side-lobe response of the telescopes as mentioned briefly and shown in Section 3.2.1) swamped the IPS signal. For the usable observations (13:15 UT – 16:20 UT), the P-Point lay 84 R⊙ off the east limb of the Sun, at a heliographic latitude of 57° north. When the observations of IPS from before 16:20 UT on 13 May
2005 were fitted using a three-mode weak-scattering model using boundaries between different regions of flow determined from the 3-D density STELab tomographic reconstruction
(as previously mentioned), they were found to be dominated by fast wind with an outflow
speed of 753 km s−1 . The precise error bound for this speed is hard to quantify in the current
three-mode model, but it lies in the region of ± 20 km s−1 . The slow-wind contribution to
94
M.M. Bisi et al.
the observation was small, and no attempt was made to fit a slow-wind speed (an outflow
speed of 350 km s−1 was assumed when fitting for fast-wind speed). Skew in the CCF peak
suggested the presence of slower fast-wind outflow near the slow-wind boundary layer (see,
e.g., Bisi et al., 2007a) and fitting indicated an outflow speed of 650 ± 50 km s−1 for this
mid-latitude component of the solar wind. J0319+415 was again observed by EISCAT and
MERLIN on 14 May 2005 from 14:00 UT, when the P-Point lay 84 R⊙ off the east limb of
the Sun at a latitude of 58° north. For the first five minutes of the observation, the results were
very similar to those from the previous day: a fast-wind speed of 760 km s−1 for the main fast
stream and ∼ 650 km s−1 for the low-latitude fast wind near the slow-wind boundary. Owing
to the short integration time, error bounds were larger than the previous day, being of the order of ± 35 km s−1 for the main fast stream and ± 75 km s−1 for the low-latitude fast wind.
From 14:05 UT to 14:10 UT, the form of the CCFs changed abruptly, with the emergence of
a third peak at shorter time lags. Considerable temporal evolution of the CCFs was obvious,
although it is difficult to determine how much of this is temporal change in solar wind parameters and how much is a result of the geometry of the observation changing with time,
maximising sensitivity to outflow at slightly different meridional angles (Moran et al., 1998;
Bisi, 2006; Breen et al., 2008).
The sudden change in the form of the CCFs and spectra indicated the onset of a transient
event in the solar wind, while the appearance of a bay of negative correlation flanking the
peak of maximum cross correlation between the scintillation patterns suggested field rotation (Klinglesmith, Grall, and Coles, 1996; Canals, 2002; Jones, 2007). Jones et al. (2007)
showed that the time of appearance of this transient was consistent with it being an interplanetary extension of the CME seen in SOHO|LASCO on the previous day. When the
observations were fitted again, by Breen et al. (2008), the results showed that the main peak
of the CCF corresponded to fast wind near the boundary layer of the ICME, deflected polewards by ∼ 8° – 15°, and with an outflow speed of 719 km s−1 . This was also consistent
with preliminary off-radial findings by Bisi (2006) for the general structures at the time. The
low-latitude fast wind appeared as a distinct peak at longer lags, with an outflow speed of
∼ 650 km s−1 , while the “new” small peak at shorter lags corresponded to the ICME and
had an outflow speed of 1075 km s−1 (Breen et al., 2008). Error bounds are of the order
of ± 15 km s−1 for the deflected fast stream, ± 40 km s−1 for the low-latitude fast stream,
and ± 100 km s−1 for the ICME. The relatively-small contribution made by the ICME to
the observed scintillation pattern is a result of much of the structure lying relatively near to
the Earth in the LOS (as mentioned previously and seen on close inspection in Figure 29).
As a result, not all the phase scintillations introduced by refraction by small-scale irregularities within the ICME had been converted to amplitude scintillation by the time the signal
reached the telescopes.
Figure 30 shows the mass determination (left panels) using the STELab IPS densityreconstruction results along with the reconstructed-velocity structure surrounding the mass
(right panels) at each of the times shown. As in Figure 29, all non-ICME associated features
have been removed. The dates and times showing the propagation of the ICME through the
interplanetary medium from near the Sun to around 1 AU are (top, going downward) 14 May
2005 at 03:00 UT and 15:00 UT (interpolated), and 15 May 2005 at 03:00 UT and 15:00 UT
(interpolated). The excess mass above an assumed background of 5 cm−3 obtained from each
of the four times during the ICME propagation are approximately constant, varying little,
with an excess mass of the ICME of around 6.4 × 1015 g, an energy fluctuating between
2.9 × 1031 erg and 4.7 × 1031 erg, encompassing a volume increasing from 0.045 AU3 to
0.204 AU3 throughout the period of interest as shown. The excess mass obtained here is
somewhat consistent with the modelled (assumed) values used by Lara and Borgazzi (2009)
The 13 May 2005 Coronal Mass Ejection
95
Figure 30 As in Figure 29 for
the most part, but this time
showing the developing and
changing reconstructed structure
of the 13 – 15 May 2005 ICME as
it makes its way to the Earth at
the times shown. The left-hand
images are those of density but
highlighted with cubes to
encompass the reconstructed
volume of the mass portion of the
ICME. This same highlighted
volume is shown on the
right-hand side displaying the
reconstructed portions of very
high velocity around the denser
ICME. Each image is labelled
with the corresponding date and
time as reconstructed along with
the masses, volume, and energy
values in each.
in their attempts to model using hydrodynamic theory the ICME transport with associated
Type II bursts. Their assumed value for a typical CME, and used for this event, is ∼ 1016 g:
around 50% more than we compute here.
An overall summary of observations of IPS from the IPS-capable systems discussed here
can be seen in Table 6. Here, the separated portions of the EISCAT – MERLIN (E – M) LOS
96
Table 6 Summary results table of observations (Obs.) of IPS for the 13 – 15 May 2005 CME/ICME/MC event. A “–” value means this is either not possible to be measured or
currently not available. Velocity values stated are the radial components of velocity as measured by each IPS system. EISCAT – MERLIN observations are shortened to “E – M”.
IPS Results Summary Organised by Date/Time
IPS Obs.
Source
Date of
Time of
Feature – Sun
Rad. Vel.
Rad. Vel.
g-level
Off-Rad.°
System(s)
Observed
Obs.
Obs. (UT)
Dist. (R⊙ )
(km s−1 )
Err. (km s−1 )
Value
(poleward)
STELab
J0323+053
13 May 2005
02:49
± 28
0.978
–
0.903
–
± 92
–
–
≈ 1.0
–
–
0°
–
–
–
02:58
≈ 48
500
J0336+322
54
564
J0534+193
04:45
115
940
Pushchino
J0137+331
12:20
> 800 (fast wind)
–
E–M
J0319+415
15:00 – 18:00
≈ 86
84 – 119
753 (fast wind)
11:00
≈ 120
750
± 20
97
870
–
–
–
≈ 50
1086
± 192
1.30
–
–
1093
± 48
1.95
≈ 60
± 20
1.23
–
5.24
–
≈ 86
> 800
± 54
–
≈ 1.3
–
760 (fast wind)
–
∼ 0°
Ooty
STELab
Pushchino
E–M
14 May 2005
J0352-071
11:16
J0323+053
02:42
J0336+322
02:54
J0431+204
03:38
J0534+193
04:42
J0137+331
12:20
J0043+520
13:20
J0319+415
15 May 2005
± 71
≈ 52
810
≈ 112
1161
83 – 129
≥ 1300
–
Before 14:10
≈ 151
After 14:10
83 – 129
719 (non-radial fast wind)
88 – 166
1075 (ICME)
137
949
–
± 15
–
∼ 8° – 15°
–
–
± 35
± 100
–
∼0°
J0158-141
J0447-220
04:03
153
910
–
–
–
Pushchino
J0137+331
12:30
≈ 86
≈ 700 (fast wind)
–
≈ 1.3
–
≈ 86
≈ 1.0
–
J0137+331
J0043+520
13:30
16 May 2005
12:30
13:30
≈ 170
≈ 150
≈ 400 (slow wind)
–
–
–
–
–
–
≈ 2.5
–
≈ 1.0
–
M.M. Bisi et al.
Ooty
J0043+520
01:14
≈ 2.5
The 13 May 2005 Coronal Mass Ejection
97
Figure 31 An overview plot of ACE solar wind measurements around the Sun – Earth L1 point. A full
description of the parameters shown can be seen in the text.
are labelled dependent on what values are quoted in the table such as the fast solar wind
or the ICME itself. The EISCAT – MERLIN values are taken from Breen et al. (2008), but
some of the error values have now been re-evaluated and shown in greater detail here.
3.2.3. Near-Earth In-Situ Measurements
Figure 31 shows the in-situ magnetic-field and plasma measurements from the ACE spacecraft just upstream of Earth. The panels, from top to bottom, show the magnetic-field magnitude, the in- and out-of-ecliptic magnetic-field angles, the proton bulk speed, the proton
density, and the proton temperature. The first solid line shows the IP shock, the second solid
line the ICME leading edge (and southward magnetic-field orientation); a possible location
of the ICME trailing edge is shown by the dashed line.
When an ICME/MC has a very high velocity (as with this ICME), it significantly compresses the solar wind plasma ahead of it and forms a “collisionless” shock, as shown
schematically in Figure 32 taken from Tsurutani et al. (2003). The shock wave (caused
by the ICME/driver gas) is formed downstream of the ICME/driver leading edge, which not
only compresses, but also heats and accelerates the ambient solar-wind plasma ahead of it.
In the sheath region between the shock and the ICME leading edge, the magnetic field “piles
up” (Siscoe and Odstrcil, 2008) while the plasma is deflected about the obstacle to the flow
(Owens and Cargill, 2004; Bisi, 2006; Breen et al., 2008).
The ICME itself is strongly characterised by a smooth rotation in the magnetic-field direction, signifying the passage of a magnetic flux rope (see, e.g., Burlaga, Behannon, and
Klein, 1987; Lepping, Burlaga, and Jones, 1990; Burlaga, Lepping, and Jones, 1990). Attempts to model the flux rope and reconstruct the large-scale topology of the ICME are
98
M.M. Bisi et al.
Figure 32 Schematic showing
the general overall configuration
of a fast CME/ICME/driver and
its downstream sheath and
subsequent shock (adopted from
Tsurutani et al., 2003).
discussed later in this section and also in Section 3.2.2 respectively. Dasso et al. (2009)
modelled the ICME as two separate MCs using two different models (both different from
the one we use here) due to the large size of the overall ICME and the possible double source
(the two flare signatures for example). The ICME also exhibits a declining velocity profile
(as discussed in Section 3.2.4), indicative of the structure expanding as it passes over the
spacecraft (see, e.g., Owens et al., 2005), consistent with the low proton temperature and
density within the flux rope.
Since fast CMEs/ICMEs generate IP shocks as they press into the slower solar wind
downstream (sticking with the previously-defined convention in Section 1 of “upstream”
being in the direction back to the Sun), using the observed photospheric magnetic field as a
boundary condition, the steady-state solar wind around the time of the 13 May 2005 CME
was reconstructed with the Enlil MHD model (see, e.g., Odstrcil and Pizzo, 2002; Odstrcil,
Pizzo, and Arge, 2005). The SOHO|LASCO white-light observations were fit with a “cone”
model of a CME (Xie, Ofman, and Lawrence, 2004), which enables determination of the
source location, time of eruption, CME angular width and radial speed from the plane-ofsky data. An over-dense plasma bubble with the same characteristics was then inserted into
the simulated ambient solar wind in order to predict the ICME-arrival time and speed at
Earth (Xie, Ofman, and Lawrence, 2004; Odstrcil, Pizzo, and Arge, 2005). The arrival time
was approximately four hours off, and the speed was approximately the same as measured
in situ (∼ 1000 km s−1 ) (Xie, Gopalswamy, and St. Cyr, 2009). This is a good result, given
the various difficulties in modelling such a complex event, but the uncertainties with this
simulation seem very large indeed. Our overall findings do not disagree with those of Xie,
Gopalswamy, and St. Cyr (2009).
Using the shock travel time from the Sun to the Earth (33 h) and its arrival velocity at
1 AU (∼ 1097 km s−1 ), Lara and Borgazzi (2009) fit the parameters for linear and turbulent
drag to the Type II burst spectrum to determine the viscosity (linear) and drag coefficient
(turbulent) of the downstream solar wind. Their model assumes the velocity varies linearly
from the Sun to the Earth allowing for the position of the CME to be determined in time. The
downstream solar wind density decreases by the factor r −2 , thus allowing the determination
of the plasma frequency at that position as a function of time. The CME driving the shock
The 13 May 2005 Coronal Mass Ejection
99
Table 7 Table showing a summary of CME and shock characteristics.
Key Characteristics of the 13 May 2005 CME and Associated Shock
Description
Value
Initial velocity near the Sun
POS 1689 km s−1
Radius stretch factor
0.77
Assumed CME mass
1016 g
Downstream solar wind velocity
415 km s−1
Viscosity (laminar)
1.03 × 1021 cm2 s−1
Drag coefficient (turbulent)
Shock in-situ (average of laminar and turbulent calculations)
3.94 × 104
∼ 1111 km s−1
is allowed to expand radially giving the cross-sectional area across which the drag forces
are acting. Table 7 shows the results of their fits. We find the drag coefficient determined
by this method ≈ 4 × 104 to be remarkable (some 4 – 5 orders of magnitude different) in
comparison to the coefficient determined by the model used in Yurchyshyn et al. (2006) of
just 1.15. This discrepancy is worth further study considering the positive outcomes of both
lines of investigation.
The onset of the event in-situ is marked by the co-occurrence of a spike in both the
density and velocity and typically is where the shock commences. These can be seen in Figure 33 using hourly-averaged data from Sun – Earth L1 in-situ instrumentation: the SOHO
– Charge, Element and Isotope Analysis System (SOHO|CELIAS) (Hovestadt et al., 1995);
the Wind – Solar Wind Experiment (Wind|SWE) (Ogilvie et al., 1995); and ACE|SWEPAM.
The end of the shock event is marked by the return of the density value near that just before
the shock event. The baseline of the density is then determined by the average of the density
just before and after the event. The normalised density peak is the maximum density that
occurs for the duration of the event with the baseline subtracted out. The value given at the
bottom of each figure is the summed particle count that passed by the spacecraft throughout
the duration of the shock. The ACE and Wind count are approximately the same, but for
some reason, the SOHO count is about 50% more than that of the other two spacecraft results. SOHO also saw a much more varied velocity structure. It is possible that the Wind and
ACE spacecraft were seeing very similar parts of the ICME/MC shock/sheath, but SOHO
was perhaps crossing through a slightly-differing region from Wind and ACE.
The baseline of the velocity is determined only by the value just before the onset of the
shock event. Unlike with density, the velocity takes much longer to reduce back to baseline
value due to the radially-extended size of the ICME – shock – sheath combination. The normalised velocity peak is the peak velocity, coinciding with the peak in density, and with the
baseline velocity subtracted out. Note however, that the peak velocity is not necessarily the
maximum velocity achieved during the entire event but at the commencement of the shock.
A summary of these values is given in Table 8.
The particles are counted using the density with the density baseline subtracted out over
the duration of the event as we have defined here. The actual velocity is used, no baseline
is subtracted out. This can be justified due to the fact that the measured velocities are much
more consistent among the different spacecraft instrumentation, in general, as compared
with the density values (which for other events has been seen to vary by factors of two to four
from instrument-to-instrument measurements), which is not quite the case here however.
Unfortunately, a comparison with the IPS using the technique of Jackson et al. (2010) with
100
M.M. Bisi et al.
Figure 33 Solar-wind
hourly-averaged velocity (dashed
line) and density (solid line)
profiles of the near-Earth
solar-wind plasma for the 14 – 15
May 2005 ICME shock/sheath
event measured by
SOHO|CELIAS (top),
Wind|SWE (middle), and
ACE|SWEPAM (bottom). The
three instruments show a density
enhancement of approximately
16.9 cm−3 to 20.7 cm−3 from
the baseline (horizontal dotted
line) and a peak velocity of
approximately 864 km s−1 to
953 km s−1 .
SMEI (and STELab IPS velocity combined) reconstruction results is not possible for this
event.
Figure 34 shows the electron (top) and proton (bottom) fluxes measured by the Electron, Proton and Alpha Monitor (EPAM) (Gold et al., 1998) aboard the ACE spacecraft.
These energetic-particle measurements clearly delineate the events previously identified in
this study. The flare onset (a) is associated with an immediate electron enhancement and a
delayed proton enhancement. Post flare, the electron and proton fluxes steadily rise, peaking
with the arrival of the shock at L1 (b). Thus, the shock front continued to act as the site of
strong electron and proton acceleration throughout its journey from the Sun to the Earth.
Particle fluxes dropped sharply with the arrival of the MC leading edge, suggesting particles
from the intense magnetic fields and closed structure associated with the magnetic flux rope.
Small enhancements of protons within the cloud, however, suggest that secondary flares may
The 13 May 2005 Coronal Mass Ejection
Table 8 Summary results table of the hourly-averaged in-situ data from SOHO|CELIAS, Wind|SWE, and ACE|SWEPAM solar-wind plasma measurements for the 14 – 15
May 2005 ICME shock/sheath.
Summary results from in-situ spacecraft measurements
Source
Start
End
Time
Density
Dens.
Norm.
Velocity
Vel.
Norm.
Particle
Date/Time
Date/Time
Span
Baseline
Peak
Den. Peak
Baseline
Peak
Vel. Peak
Count
(UT)
(UT)
(h)
(# cm−3 )
(# cm−3 )
(# cm−3 )
(km s−1 )
(km s−1 )
(km s−1 )
(#)
SOHO
15 May 01:00
15 May 11:00
10
2.93
22.64
19.71
398.9
864.1
465.2
Wind
14 May 22:00
15 May 09:00
11
2.25
19.10
16.85
414.0
952.7
538.7
ACE
14 May 21:00
15 May 10:00
13
2.71
23.44
20.73
415.1
946.3
531.2
3.14 × 1013
1.92 × 1013
2.19 × 1013
101
102
M.M. Bisi et al.
Figure 34 Electron (top) and proton (bottom) fluxes measured by the ACE|EPAM instrument. The solid
vertical lines show (a) the EIT flare onset, (b) the observed shock arrival at L1 and (c) the MC leading edge.
The dashed vertical lines show the MC trailing edge, as inferred from (d) the end time of “ideal” in-situ
magnetic-field measurements (see later in Section 3.2.4 of the text for explanation), and (e) the force-free
model.
have injected particles into the ICME itself (see, e.g., Larson et al., 1997). Finally, we note
that, as with other interplanetary signatures, the trailing edge of the MC is more difficult
to identify: lines (d) and (e) show the “ideal” in-situ magnetic-field measurements end time
(see Section 3.2.4) and the force-free predictions, respectively.
3.2.4. A Brief Overview of the Inner-Heliosphere Observations and Measurements
Table 9 combines the shock speed values obtained from the DH-km-TII analysis with coronagraph white-light observations, observations of IPS, and in-situ measurements, in order
to study the temporal speed evolution of the CME/ICME/shock. Column 1 shows the date
and Column 2 the time of the measurement as a fraction in UT. Column 3 shows the speed
along with its uncertainty in Column 4. Column 5 identifies the instrument/technique used to
determine the speed measurement. Since the four IPS-capable systems (EISCAT, MERLIN,
STELab, and Ooty) have different receiver/antenna types, the determination of radial velocity (Vips ) is undertaken using different methods. The ICME was detected in-situ by the Wind
spacecraft. The local shock speed (Vshock ) inferred from Wind in-situ data by employing the
Rankine – Hugoniot relations. The last two rows of the Table show the shock (Vshock ) and
ICME transit speed (Vicme ), from taking the LASCO – CME observations as the initial time
and the shock and Wind – ICME observations as the final time. These values show a clear
deceleration of both disturbances: the transit speeds are faster than the local ones (Wind
in-situ).
The 13 May 2005 Coronal Mass Ejection
103
Table 9 Results of speed determinations by different instruments of the 13 – 15 May 2005
CME/ICME/shock. EISCAT – MERLIN observations are noted as “E – M” and Ooty results are an overall
summary of multiple observations. When an “a” is marked next to the date, this denotes when the solar wind
was sampled and not that of the CME/ICME/shock; a “–” marks where these values are not available/not
measured.
13 – 15 May 2005 CME/ICME/Shock Speeds
Date
Hour
Speed
Error in Speed
Instrument or
(2005)
(UT Fraction)
(km s−1 )
(± km s−1 )
Technique used
13 May
17.2
(POS) 1689
–
LASCO (CME)
13 May
18.2367
1176
40
13 May
22.5417
1230
30
WAVES – RAD1
13 Maya
13.6667 – 16.3333
753
20
E – M IPS (high-lat. wind)
WAVES – RAD1
13 Maya
13.6667 – 16.3333
650
50
14 May
2.5
1050
150
14 May
2.7
1086
192
14 May
2.9
810
48
STELab IPS
Wind – RAD1
E – M IPS (mid-lat. wind)
Ooty IPS (ICME)
STELab IPS (ICME)
14 May
3.75862
826
206
14 May
3.8
1093
20
14 May
4.7
1161
54
14 May
12.0
1050
150
14 May
12.15
916
92
WAVES – RAD1
14 Maya
14.0 – 14.0833
760
35
E – M IPS (high-lat. wind)
14 Maya
14.0 – 14.0833
650
75
E – M IPS (low-lat. wind)
14 Maya
14.1667 – 18.0
719
15
E – M IPS (high-lat. wind)
14 Maya
14.1667 – 18.0
650
40
E – M IPS (low-lat. wind)
14 May
14.1667 – 18.0
1075
100
E – M IPS (ICME)
15 May
2.25
1100
–
Wind|SWE (shock)
15 May
6.0
950
–
Wind|SWE (ICME)
15 May
6.0
900
100
Ooty IPS (ICME)
100
STELab IPS (ICME)
STELab IPS (ICME)
Ooty IPS (ICME)
15 May
17.0
900
15 May
2.25
1261
0
transit speed (shock)
Ooty IPS (ICME)
15 May
6.0
1132
0
transit speed (ICME)
The interplanetary velocity changes through the progression of the 13 – 15 May 2005
CME/ICME/shock can be seen in Figure 35. This figure shows the interplanetary velocity
evolution of the ICME/shock associated with the 13 May 2005 CME. The combination of
all the CME/ICME/shock speed determinations show a gradual deceleration as it propagates
toward 1 AU, but probably not overall linearly close-in to the Sun.
This comparison of speed determinations should be done carefully. In general, we assume
that the Type II radio emission is coming from the region around the ICME/shock nose, but
sometimes it may come from other regions along the ICME/shock front. The near-Sun POS
LASCO-CME speeds, the shock speeds from the km-TII slope analysis, the sheath region
speeds from IPS observations, and the Wind-shock and ejecta speeds from in-situ magnetic
and plasma data do not refer to the same structure but are related to the same phenomena
(CME/ICME/shock event).
The Type II radio-burst tracking technique can be summarised as follows. Intervals where
the DH-km-TII emission was not contaminated by another kind of emission (i.e. Type III
104
M.M. Bisi et al.
Figure 35 Speed temporal heliocentric evolution of the 13 May 2005 LASCO CME. Combination of speed
determinations by different instruments: white-light SOHO|LASCO (open triangle); shock speed obtained
from the analysis of the radio km Type II burst drifts (solid circle); IPS STEL/STELab (open diamond); IPS
ORT/Ooty (solid diamond); IPS EISCAT-MERLIN/E-M (solid square); and in-situ Wind plasma measurements (open square).
radio bursts which occur during this flare/CME event, auroral kilometric radiation – AKR,
etc.) were identified. The DH-km-TII emission was then isolated and a Gaussian fit (for all
spectral data) to the profile of the flux density versus frequency applied to obtain its spectral properties (central frequency, width, and intensity). We only considered those spectra
with a well-defined central frequency. Next, the slope of the central frequency drift was obtained by applying a linear fit at some arbitrary interval. Using Equation (3) and the slope
of the central frequency emission, the shock speed was then estimated and then these steps
repeated at different contiguous intervals to obtain a shock speed distribution. Finally, the
mean value of the shock speeds is found and its standard deviation quoted as the associated
error. With comparison to the other velocity determinations from near the Sun through to
near Earth, we see that there is a strong deceleration curve (not linear near the Sun – suggesting rapid deceleration close to the Sun – but possibly linear deceleration in interplanetary
space and out to the Earth) and that our various data sets compare reasonably well with this
picture (as shown in Figure 35). There is some discrepancy and question over the originallyquoted-and-regularly-used SOHO|LASCO POS velocity of 1689 km s−1 determination in
comparison with our much lower bounds (for the two separate fronts before their apparent
“merging”) using the new technique as seen in Section 3.1 specifically relating to Figure 14.
The overall STELab results show an enhanced turbulence and velocity in the vicinity
of the ICME as it propagates out from the Sun. Although the enhanced g-level value (turbulence/density increase) was inconclusive, the velocity distribution shown by the STELab
The 13 May 2005 Coronal Mass Ejection
105
observations of IPS consistently show an Earth-directed halo-type transient, but with some
bias to the north and east. These are somewhat consistent with the SOHO|LASCO analyses.
The overall Ooty results show the passage of the ICME very well on the 14 and 15 May
2005. The observations are also widely spaced in terms of solar elongation (and thus P-Point
distance from the Sun). The ICME crosses the distance of around 100 R⊙ to 150 R⊙ on the
14 May with velocities in the range of 750 km s−1 to 1150 km s−1 . By 15 May, a halo-like
ICME can be seen in the Ooty data, which seems to be primarily moving around 10° – 20°
east of the Sun – Earth line despite a greater number of IPS radio sources being observed
to the west (as with the largest of the mass portions seen in the SOHO|LASCO images
and in the STELab 3-D reconstructions). The ICME appears to start to cross the Earth’s
orbit (solar elongation of 90°) around 06:00 UT – 12:00 UT on 15 May. This is consistent
with the timing of near-Earth in-situ measurements inside the Earth’s magnetosphere for
the ICME/MC and the shock/sheath arrival/impact. On 16 May 2005, the ICME passes
250 R⊙ and moves away from the Ooty IPS FOV. However, late on 15 May 2005, there are
some ICME signatures around 225 R⊙ . These observations are in good agreement with the
measurements made by the Wind spacecraft instrumentation in the interplanetary medium
near Earth.
The ability of ELB IPS observations to detect meridional velocities (see, e.g., Moran
et al., 1998; Bisi et al., 2007a) revealed significant polewards deflection of the fast wind
near the ICME boundary north of the ICME. This is consistent with over-expansion of the
ICME into the fast wind, offering a possible explanation for the north – south feature and
lower drag coefficient suggested by Yurchyshyn et al. (2006).
The excess mass, as determined from the 3-D tomographic reconstruction of STELab
IPS observations, show a fairly consistent and stable value in the region of 6.4 × 1015 g,
with energy values between 2.9 × 1031 erg and 4.7 × 1031 erg. The volume increased during its interplanetary transit (as shown in Figure 30) from 0.045 AU3 to 0.204 AU3 . Lara
and Borgazzi (2009) assumed a value for a “typical” CME, when modelling this event, of
∼ 1016 g: around 50% more than we compute here as an excess mass. However, it should be
noted that we are measuring only the excess mass above the ambient. If the ambient were
to be added as a total mass for modelling purposes, then the mass would be (our computed
total mass values) around 1.4 × 1016 g; this is about 40% more than what was used as an
assumption by Lara and Borgazzi (2009) for their modelling. It is likely that the actual mass
of the ICME is indeed ∼ 1016 g, which is consistent with what we calculate here and what
was used during the Lara and Borgazzi (2009) modelling.
The 13 May 2005 CME had a smoothly-varying magnetic field, a low proton temperature, and a well-defined counter-streaming-electron flow, indicating that magnetic-field lines
in its structure remained connected to the photosphere (Gosling et al., 2006). The orientation
of the magnetic flux rope within the ICME was roughly the same at 1 AU as the inferred orientation of the CME based on the density enhancements from its legs in the SOHO|LASCO
measurements (Yurchyshyn et al., 2006). However, there is the likelihood of the CME having a double source due to the two flares quite close together in time (the more powerful
of the two being the second flare) and also a possibility of there being two MCs within the
ICME as it passes over the Sun – Earth L1 point (as modelled by using two different models, one for each of the two suspected MCs, by Dasso et al., 2009). Table 10 summarises
the parameters used to fit the ACE data here using the Lepping, Burlaga, and Jones (1990)
magnetic-flux-rope model (see http://lepmfi.gsfc.nasa.gov/mfi/mag_cloud_pub1.html).
Plots of the fit to the flux rope detected in ACE in-situ measurements using the method of
Lepping, Burlaga, and Jones (1990) can be seen in Figure 36. The Lepping fit, shown with
the blue curves, indicates that a force-free solution with no expansion of the rope as it passes
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Table 10 Table showing a summary of ACE ICME fitted parameters.
Key Fitted Parameters from ACE ICME Measurements
Description
Force-Free
Ideal
Start time
15 May 2005 05:42 UT
15 May 2005 05:42 UT
End time
15 May 2005 22:18 UT
15 May 2005 22:18 UT
Axis direction (GSE) longitude and latitude
94° and 67°
90° and 0°
Average solar wind speed
843 km s−1
–
Radius
0.195 AU
–
Axial field magnitude
70.6 nT
70.6 nT
Handedness
Left
Left
Square-root of chi-squared of fit
0.123
–
Asymmetry factor
0.5%
–
CA (of axis relative to the Earth)
75%
80%
Axial magnetic flux
82.5 Mx
82.5 Mx
Axial current density
4.6 µA km2
4.6 µA km2
Cone angle
92°
90°
Quality of fit (1 = best; 3 = bad)
2
–
the Earth is too simple for the 15 May 2005 ICME. The fit illustrates two characteristics that
indicate the orientation is in error: the magnitude of the total field decreases throughout the
time interval, and the strong rotation in the Bz magnetic field from negative to positive is
not included in the fit. The Bz rotation indicates that the orientation of the axis is at a lower
latitude. ACE data are also shown, and are in black. To test this theory, we set the trailing
edge of the structure at the point where the velocity fields in the GSE Y and Z directions
dropped to zero; this indicated a change in the dynamics within the ICME structure. The
red curves in the figure show the ideal in-situ magnetic-field measurements that would be
observed by a flux-rope oriented along the GSE Y -axis with an impact parameter at 80%
of the radius and a 65 nT magnetic-field magnitude. Clearly the ICME is oriented in the
east-west direction; therefore, we conclude that the Yurchyshyn et al. (2006) white-light
density enhancement in the north-south direction is uncorrelated with the orientation of the
ICME. This may be consistent with the two MC structure described by Dasso et al. (2009)
within the ICME near Earth along with the bias in density to the north and east as seen in
the post-processed LASCO imagery and the overall IPS and 3-D reconstruction results.
Also of interest are the two highly-different drag coefficients used by Yurchyshyn et al.
(2006) and Lara and Borgazzi (2009) being 1.15 and ∼ 4 × 104 respectively, but still resulting in similar and perhaps suitable results as compared with near-Earth in-situ data. As
mentioned in Section 3.2.3, this certainly merits further investigation.
3.3. The Earth and its Environment
Here, we describe the effects on the Earth environment and study in detail the magnetospheric response to the 13 – 15 May 2005 CME/ICME/MC event. During the interval studied in the proximity of Earth (15 May 2005 – 00:00 – 12:00 UT), the Geotail, four Cluster,
and Double Star TC2 spacecraft were all initially located in the dayside magnetosphere and
were thus ideally placed to observe the effects of the impact of the ICME/MC, and the shock
and sheath preceding it on the dayside magnetosphere. Unfortunately, there were no spacecraft located in the magnetotail during this interval. Spacecraft locations plotted in the GSE
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Figure 36 The Lepping fit (dark blue curve) to the ICME ACE in-situ measurements (black) along with an
ideal curve (in red) that would be measured if the flux rope were oriented along the GSE Y -axis. Full details
can be found in the text.
coordinate system (where X points along the Earth – Sun line, positive towards the Sun; Y
opposes the orbital motion of the Earth, positive towards dusk; and Z points out of the plane
of the ecliptic, positive northwards), can be seen in Figure 37, where distances are in R⊕ .
3.3.1. Multi-spacecraft Measurements and Ground-Based Observations of the Effects on
the Earth’S Magnetosphere and Ionosphere
At 00:00 UT on 15 May the Geotail spacecraft was located on the dawnward flank of
the magnetosphere, moving towards perigee close to the location of a model magnetopause. Magnetic-field data from Geotail’s Magnetic Field Investigation (MFI) instrument
(Kokubun et al., 1994) and ion moments from the Low-Energy Particle (LEP) experiment
instrument (Mukai et al., 1994) are plotted in Figure 38. Where multiple vector components
are plotted in one panel, the X component is red, Y green, and Z blue. At 02:40 UT (marked
by the first dashed line) the spacecraft crosses the magnetopause into the magnetosheath (as
evidenced by a sharp increase in the magnetic-field strength, ion density and ion velocity,
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Figure 37 Geotail (green) is
skimming the dayside
magnetosphere, close to the
location of a model
magnetopause. Cluster (blue) is
outbound from the
magnetosphere into the dawnside
magnetosheath. LANL-02A (red)
is in a geosynchronous orbit.
Spacecraft locations at 00:00 UT
are marked by pluses, spacecraft
locations at 02:40 UT are marked
by solid squares.
and decrease in ion temperature). This turbulent sheath plasma and magnetic field were measured by the spacecraft until ∼ 06:00 UT (marked by the second dashed line), when the Z
component of the magnetic field abruptly turned southward and the field and ion velocity
became much less variable. Note that spacecraft data are plotted in GSM coordinates, which
are given by a rotation of GSE about the X-axis such that the GSM X – Z plane contains the
Earth’s magnetic-dipole axis, rather than the normal to the plane of the ecliptic.
The orbit of the Cluster spacecraft (only data from Cluster 1 will be discussed here) was
such that the spacecraft crossed the magnetopause and moved from the magnetosphere into
a “normal” undisturbed magnetosheath at 02:00 UT before Geotail moved into the magnetosheath. Several Cluster instruments were used for this investigation. The magnetic-field
data from the Flux Gate Magnetometer (FGM) (Balogh et al., 2001), ion moments from the
Cluster Ion Spectrometry – Hot Ion Analyser (CIS – HIA) (Reme et al., 2001), and an electron energy-time spectrogram from the Plasma Electron And Current Experiment (PEACE)
(Johnstone et al., 1997), are plotted in Figure 39. The 02:00 UT crossing of the magnetopause by the Cluster 1 spacecraft, marked by the first dashed line, can be seen as the
change from higher energetic magnetospheric electrons to colder magnetosheath electrons
of higher flux at that time. At 02:40 UT (marked by the second dashed line) however, the
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Figure 38 From the top, the panels show magnitude and GSM components of magnetic field, GSM components of v⊥ , ion number density, and ion temperature, as measured by the Geotail spacecraft.
spacecraft measured a change in the electron distribution, accompanied by an increase in the
ion velocity and slight increase in the ion temperature. This hotter, faster sheath plasma continued to be observed until ∼ 06:00 UT (marked by the third dashed line) when, like Geotail,
the Cluster spacecraft measured a change in BZ to a southward direction, accompanied by
a decrease in electron fluxes and temperatures.
The change in BZ to a southward direction (i.e. a negative value of θ ) was also detected
by the Wind spacecraft, and by the ACE spacecraft at ∼ 05:30 UT at the transition between
the sheath of the ICME and the MC flux-rope structure (see Figure 31 showing the ACE data,
third panel down, second vertical line). It is our interpretation of the Geotail and Cluster 1
data therefore, that the IP shock associated with the leading edge of the ICME impacted
the Earth’s magnetosphere at 02:40 UT, compressing the dayside magnetopause and hence
causing Geotail to move from the magnetosphere to the magnetosheath. The change in the
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Figure 39 From the top, the panels show magnitude and GSM components of magnetic field, GSM components of ion velocity, ion temperature parallel and perpendicular to the magnetic field, and an energy – time
spectrogram plotting omnidirectional differential energy flux of electrons. Dashed lines mark the crossing
of the spacecraft from the magnetosphere into the magnetosheath, from undisturbed magnetosheath into the
sheath of the MC and finally into the MC itself.
plasma measured by Cluster 1 at this time from “normal” magnetosheath to hotter plasma
is also consistent with the leading edge of the ICME sheath crossing that spacecraft at this
time.
If the Geotail data are plotted with various geomagnetic indices (sym-H, AU, AL), an
indication of the effect of the IP shock and MC on the magnetosphere – ionosphere system
can be ascertained (see Figure 40). First considering sym-H, which is broadly equivalent to
Dst but at a one-minute resolution (see, e.g., Wanliss and Showalter, 2006), a sharp increase
in sym-H can be observed at 02:40 UT, when the IP shock impacted the magnetosphere and
Geotail was pushed into the magnetosheath. The increase in sym-H is consistent with the
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Figure 40 From the top, the panels show the sym-H index (black) along with the Dst index (red), Kp, PCN,
AU and AL indices, the magnitude of magnetic field as measured by Geotail, GSM BZ from Geotail, and
GSM v⊥ X also from Geotail.
initial phase of a geomagnetic storm and is caused by an increase in the magnitude of the
magnetopause current resulting from the increase in solar wind dynamic pressure (Gonzalez
et al., 1994). The IMF in the ICME sheath was predominantly directed northward (Figure 31
third panel) so one might not expect increased dayside reconnection rates during the initial
phase of the storm. The main phase of the geomagnetic storm, identified as the decrease in
sym-H, began at ∼ 06:00 UT and was coincident with the southward turning of the IMF seen
by Geotail (Figures 38 and 40) and upstream by ACE some thirty minutes earlier (Figure 31)
as the MC itself impacted the magnetosphere. The recovery phase of the storm, identified
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as the increase sym-H to quiet-time values, began at ∼ 08:00 UT and continued throughout
the day.
The AU and AL indices (Davis and Sugiura, 1966), a measure of the strength of the
AEs, showed a transient (∼ 1 h) increase beginning at 02:40 UT; consistent with substorm
activity being triggered by the pressure pulse associated with the ICME sheath (see, e.g.,
Keika et al., 2009). The PCN index also showed a sharp increase from 5.0 nT at 02:30 UT
to 11.1 nT at 03:15 UT. It fell to ∼ 6.0 nT between 03:30 UT and 05:45 UT before a second
peak value of 10.9 nT occurred at 06:15 UT. Subsequently, it fell to ∼ 3.0 nT at 12:00 UT;
although there was significant variation within this decline. The enhanced values between
∼ 03:00 UT and ∼ 08:00 UT suggest enhanced anti-sunward convection within an expanded
polar cap. The Kp index showed a maximum value of 83 at 06:00 UT before falling to lower
values around 40 at 12:00 UT. The Dst index fell sharply around 06:00 UT, took a minimum
value of −263 nT at 08:00 UT, and then began to recover over several hours. These both
indicated geomagnetic disturbances on a global scale occurring a few hours later than the
disturbances at high latitudes.
The geomagnetic indices suggest that the ionospheric response to the ICME/MC was
consistent with that previously reported by Hanuise et al. (2006) on the May 2003 CME
events and the response to an associated storm (Ebihara et al., 2005b, 2005a). The disturbances in the AU, AL, and PCN indices were coincident with the impact of the ICME/MC
at Earth and indicated the response of the high-latitude ionosphere; namely enhancements
in the AEs and the anti-sunward cross-polar transport of plasma. Disturbances in the Dst
and Kp indices occurred some hours later, indicating a delay in the global response to the
ICME/MC.
Unfortunately, no auroral data are available until after the geomagnetic storm had begun.
Global auroral images from IMAGE Far Ultraviolet Imager – Wideband Imaging Camera
(FUV – WIC) (Mende et al., 2000) are available from ∼ 04:50 UT. These images, representative samples of which are plotted in Figure 41, show enhanced auroral brightness visible
above the background airglow on the dayside beginning as soon as the data become available and continuing throughout the interval. Unfortunately, the nightside aurora are not in
the FOV of the WIC camera until 06:00 UT. However, data from the NORSTAR MSP at the
Gillam observation site in Canada (Figure 42), which measures auroral emission at several
wavelengths over a narrow longitude range but a wide range of latitude, are available from
∼ 03:30 UT and show the equatorward edge of the nightside aurora moving to lower latitudes. This is consistent with the expansion of the ionospheric polar cap caused by enhanced
reconnection at the dayside magnetosphere producing more open magnetic flux (Milan et al.,
2006). This is interesting because the predominantly northward-directed IMF in the ICME
sheath would not usually be expected to produce such significant open magnetic flux. Component reconnection involving the Y -component of the IMF may be responsible for this;
however, recent work by Milan et al. (2009) also suggests that the strengthening of the ring
current associated with the geomagnetic storm also has the effect of expanding the radius of
the auroral oval.
Both the IMAGE FUV and Gillam MSP data show enhanced auroral activity on both the
dayside and nightside, where the aurora were much more active than for a typical auroral
substorm, as the main phase of the geomagnetic storm began at 06:00 UT. However, the
dayside reconnection rate was significantly higher than the nightside, as evidenced by a
further expansion of the ionospheric polar cap, visible in the FUV data, as more and more
open magnetic flux was added to the magnetosphere. This is entirely consistent with the
southward turning of the IMF detected by both Geotail and ACE (as previously mentioned)
at the transition between the ICME sheath and the MC itself, providing further evidence that
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Figure 41 Selected southern hemisphere global auroral images taken by the FUV-WIC camera on board the
IMAGE satellite. The images are in a geomagnetic latitude/longitude coordinate system and are plotted such
that noon is at the top of each image and midnight at the bottom.
the main phase of this geomagnetic storm was associated with the MC, and the initial phase
with compression of the magnetosphere caused by the impact of the IP shock and ICME
sheath.
The response of the high-latitude convection pattern and the Birkeland current system in
the terrestrial ionosphere to this ICME has been discussed by Eriksson et al. (2008). The
plasma convection, the Birkeland current system, and the auroral flux were observed by the
IMAGE, Thermosphere Ionosphere Mesosphere Energetics and Dynamics (TIMED) (see,
e.g., Yee, Cameron, and Kusnierkiewicz, 1999), and US Defense Meteorological Satellite
Program (DMSP) satellites. Taken collectively, these showed that a field-aligned current
system developed at the poleward edge of the dusk auroral oval when the IMF was directed
southward. When the IMF turned northward, a flow channel of sunward-directed plasma in
the northern hemisphere was observed between upward and downward field-aligned currents
and this structure drifted dawnward across several hours. The upward field-aligned current
was coincident with a drifting transpolar arc in the northern hemisphere, but there was no
counterpart observed in the southern hemisphere.
The five Los Alamos geosynchronous satellites also detected effects of the impact of the
ICME/MC on the magnetosphere. At 02:40 UT, an increase in energetic electron fluxes (not
shown here) was detected, consistent with adiabatic heating of the existing electron popula-
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Figure 42 Keograms of auroral emission in three wavelengths as measured by the Gillam NORSTAR MSP.
This figure shows the temporal evolution of the latitude of the auroral oval along a narrow longitude range.
tion caused by the compression of the magnetosphere, as was a new injection of energetic
particles in the midnight sector. This is consistent with the substorm activity triggered by the
magnetospheric compression. By 06:00 UT, the closest LANL spacecraft to local noon was
LANL-02A. This spacecraft crossed into the magnetosheath at ∼ 05:50 UT as evidenced
by the change in the thermal and suprathermal electron and ion spectra measured by the
Magnetosphere Plasma Analyzer (MPA) instruments (Bame et al., 1993) and plotted in Figure 43 (where the first dashed line marks the transition between the magnetosphere and
magnetosheath), from a magnetospheric population to a sheath population. The effect of
the ICME/MC then, was to compress the dayside magnetosphere to within geosynchronous
orbit. This of course may likely have affected nearby geosynchronous satellites.
Double Star TC2 was moving towards its apogee of ∼ 6 R⊕ . The spacecraft recorded increased electron fluxes and a dipolarisation (i.e. an increase in BZ ) at 02:40 UT. At 05:30 UT,
the spacecraft recorded a more variable, turbulent magnetic field, and a change in the electron population to a lower energy distribution dominated by field-aligned and anti-fieldaligned particles. At 06:00 UT, the spacecraft observed a southward turning of BZ , presumably related to the MC passing over the spacecraft. At this time, the bi-directional elec-
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Figure 43 Energy – Time spectrograms plotting omnidirectional differential energy flux as measured by the
MPA instruments on LANL-02. This confirms that the spacecraft drops into the sheath at ∼ 06:00 UT.
tron population disappeared and a slightly lower energy isotropic distribution was recorded.
These interesting data require further analysis before they can be properly interpreted in the
context of the ICME and MC and will likely become subject of a forthcoming paper.
3.3.2. A Brief Overview of the Earth and its Environment Observations and Measurements
The magnetospheric response was large and widespread with the magnetosphere being compressed into geosynchronous orbit. The initial phase of the compression was caused by the
impact of the IP shock and ICME sheath, but the main phase of the geomagnetic storm was
caused by the southward-turning IMF (and hence associated with the MC features). The
ICME/MC effects measured at the Sun – Earth L1 point quickly progressed onto affecting
the Earth’s local magnetic environment. A good connection has been made between features seen first measured in-situ around L1 and those a little later seen at the Earth. This
was the largest geomagnetic storm of 2005 and its effects were seen throughout the Earth’s
magnetosphere – ionosphere system.
Enhancements in the AEs, anti-sunward cross-polar plasma transport, and large geomagnetic disturbances all occurred as a result of this storm. The geomagnetic indices suggest
that the ionospheric response to this ICME/MC was consistent with that previously reported
on the May 2003 CME events by Hanuise et al. (2006) and that this also resulted in a similar
response to an associated storm (Ebihara et al., 2005b, 2005a).
The detailed response of the high-latitude plasma convection pattern and the Birkeland
current system have been discussed by Eriksson et al. (2008). When the IMF was directed
southwards, a field-aligned current system developed at the poleward edge of the dusk auroral oval. Once the IMF turned northward, a flow channel of sunward-directed plasma in the
northern hemisphere was observed between upward and downward field-aligned currents,
and over several hours, this flow channel drifted dawnward.
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4. Summary and Conclusions
The results discussed here reveal the complexity of this event, but also point towards
some conclusions – some specific to this event – and some of more general relevance to
CME/ICME/MC studies – as well as certain outstanding questions to be addressed in future
studies.
The object of this paper was to draw together as much data as possible concerning the
initiation of the 2005 solar eruption, the development of the mass ejection, its propagation
through the interplanetary medium and its impact on the Earth, and to attempt to draw some
preliminary conclusions about the development of the event. We believe that we have not
only succeeded in these aims, but have been able to go beyond them in terms of drawing
some general conclusion about the evolution of solar storms and their planetary impact. We
also believe that the results described in this paper point towards some important questions
to be addressed in future studies.
In relation to the EUV and AR magnetic-field modelling, based on these encouraging
results from the vector-magnetic field study on this event, we are now preparing to repeat this
study with the full MHD model. The SOHO|CDS intensity measurements show clear signs
of the larger flare eruption and also some signs of the earlier, weaker flare (which could also
have played a role in the early development and structure of this flare/CME eruption). Using
the new method for looking at the SOHO|LASCO coronagraph imagery, we detect a much
lower minimum velocity (POS) for a possible two separate and “merging” CME fronts,
but further study on this aspect of the event is needed before any significant conclusions
can be drawn. It does indicate however, that this new post-processing method by Morgan,
Habbal, and Lugaz (2009) could become very useful for other CMEs when a more detailed
investigation is warranted.
This paper significantly highlights the importance of observations of interplanetary scintillation (IPS). They show us the grand-scale propagation of the ICME through the interplanetary medium, its changing structure and distribution in the heliosphere, and also gives
us insight into the relatively-large non-radial poleward flows of the fast wind to the north as
a result of the ICME’s passage allowing us to look at the interaction between a travelling
ICME and the background solar wind more closely than any other method currently available. No other technique is able to do this on such a scale except for white-light observations
using the SMEI instrument (Eyles et al., 2003; Jackson et al., 2004), which unfortunately
was not operational during the time of the 13 – 15 May 2005 CME/ICME/MC due to technical problems. Other heliospheric-type imagers could possibly provide improved large-scale
results provided an adequate FOV were available (for example, at least the sunward hemisphere of the inner heliosphere in its entirety).
The post-processed LASCO coronagraph images, observations of IPS, 3-D density reconstruction of STELab IPS data, and MC flux-rope fits all show a halo-like structure coming toward the Earth, but all with a bias to the north-east quadrant. This gives some strength
to the possibility that there was a secondary eruption which quickly merged with the first in
the low corona – this is possibly why both a near-symmetric halo-type structure was seen
in coronagraph imagery as well as a north-east dominance later – particularly through the
interplanetary environment. This is also a possible agreement with the two separate MCs
detected and discussed in some detail by Dasso et al. (2009) in the interplanetary medium
near the Earth.
We suggest that one of the most important results to emerge in this work is the nonuniform nature of the event, both in the initial CME through the corona and in the ICME
in interplanetary space. The ELB IPS results show that the ICME was perturbing the
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background fast wind, driving a northward flow, and this is strongly suggestive of ICME
over-expansion into the fast wind. This would provide an explanation for the variations in
ICME velocities inferred from different observations in interplanetary space, in particular
the higher velocity seen in the relatively-high-latitude ELB IPS observation compared with
some at lower latitudes. The relative importance of initial velocity structure and subsequent
interaction with the solar wind in determining the morphology of the ICME at Earth orbit is
an important question which has emerged during this study, and we consider this to be one
of the most important contributions of this work to the overall field of Sun – Planet connections. This is (to our knowledge) the most comprehensive speed analysis of an ICME-shock
event. We are combining speed determinations by six different instruments illuminating
some aspects of the interplanetary evolution of the ICME-shock system.
The magnetospheric effects in summary, along with Cluster, Geotail, and geomagnetic
indices, all give the time of impact of the IP shock and MC sheath on the magnetosphere as
02:40 UT; while the MC itself impacted the magnetosphere at ∼ 06:00 UT. The initial phase
of the geomagnetic storm was caused by the dynamic pressure increase associated with the
ICME sheath, while the storm’s main phase was caused by the MC itself. The impact of the
ICME/MC compressed the dayside magnetopause in such a way that it was located within
geosynchronous orbit, and enhanced reconnection at the dayside caused massive substorm
activity seen in ground- and space-based auroral data.
The complexity of even this “typical” CME means that its “evolution” and effects can
only be understood by drawing on the widest possible range of data sources and analytic
methods (as we have attempted here). It is apparent that the 3-D structure of this event is
complex, and that this structure may be determined by asymmetries in the initial eruption
as well as by interaction between the ICME/MC and the background solar wind during
interplanetary transit. The 3-D structure of the ICME will also play an important role in
governing the way in which it couples into the magnetosphere and ionosphere of the Earth.
A promising direction for future work will be a comparison of the observational results of
this study with 3-D numerical modelling of CME/ICME evolution with coupling into the
terrestrial environment.
The AR eruption launched a strong impulsive event, characterised by a clear coronal radio burst and the launch of a fast CME. The SOHO|CDS results show that the initial CME
was significantly non-uniform in character, with considerable variation in velocity across it.
The interplanetary observations show that the ICME underwent considerable deceleration –
almost certainly by interaction with the slower background solar wind ahead of it – and that
most of this deceleration took place inside 80 R⊙ – 100 R⊙ . There were considerable variations in the ICME velocities measured in interplanetary space, and these appear to reflect
the non-uniform nature of the event but give overall good comparison with each other over
increasing radial distance from the Sun – the ICME was expanding more rapidly in some
regions than others, however. Determining its relative importance and determining the initial
CME structure and its subsequent interaction with solar-wind structures and determining the
interplanetary structure of the ICME were not possible from this initial study, but this study
has emerged as an important question for future work. The ICME passed over the L1 spacecraft at 02:11 UT on 15 May 2005, and the first magnetospheric response was seen only
29 minutes later – a rapid response. The ionosphere began to respond at 03:15 UT, when
the PCN maximum occurred. The event stands as a good example of a fast Earth-directed
CME/ICME/MC and thus is a good candidate for extended study.
This event (CME – ICME – MC) was originally thought to be a fairly “simple” solar eruption, and one which could be taken to be a “typical” CME/ICME/MC. The picture which has
emerged from this study is of a much more complex event, with a significantly asymmetric
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structure which, in turn, is probably important in governing its effects on the space environment of Earth. We are of the opinion that other “apparently simple” eruptions, investigated
in the same detail as has been in our study here, would reveal comparable levels of complexity, and thus that this event may be taken to be a “typical” CME in terms of its complex
and developing structure (if not in its detailed morphology). The results of this study suggest some additional specific lines of future research. It is clear that both the initial eruptive
event and the interplanetary counterpart are significantly asymmetric in structure, and that
at least part of the asymmetry apparent in the ICME is as a result of interaction between the
ejection and the background solar wind. The relative importance of the initial CME structure
and subsequent interaction in determining the ICME structure at 1 AU is, however, unclear,
as is the role played by this structure in MC/magnetosphere coupling. Comparisons of observational results with 3-D MHD modelling could prove instructive in both cases, while
future event studies – particularly for events with observations/measurements carried out
using SMEI and/or the Solar TErrestrial RElations Observatory (STEREO) (Kaiser, 2005;
Kaiser et al., 2008) twin spacecraft – could provide greater detail on the ICME structure both
in interplanetary space and in the vicinity of the Earth, particularly with higher-resolution
3-D reconstruction using SMEI brightness data.
This leaves the necessity of still more study of this event with some aspects and questions
of it (and others) left to be answered.
Acknowledgements The University of California, San Diego (UCSD) authors acknowledge the US
NSF (grants ATM0331513, ATM0852246, and ATM0925023), US Air Force Office of Scientific Research
(AFOSR) (grant FA9550-06-1-0107), and NASA (grant NNG05GG45G) for their funding while working on
these analyses and writing up this manuscript. Four of us (R.A. Fallows, A.P. Walsh, A.G. Wood, and A.S.
Giunta) were supported by the Science and Technologies Facilities Council (STFC) in the UK during the
period when this work was undertaken. Work at Predictive Science was supported by NASA (Heliophysics
Theory and Living With A Star programs) and NSF (CISM and SHINE programs). H. Morgan acknowledges the support of NASA (grant NNX08AJ07G) to the Institute for Astronomy, University of Hawaii, for
the work on CME image processing. The STELab IPS observations were carried out under the solar wind
program of the Solar-Terrestrial Environment Laboratory (STEL) of Nagoya University. The work at Ooty
is partially supported by the CAWSES-India Program, which is sponsored by ISRO. Pushchino authors acknowledge the support of the Russian Academy of Sciences Program “Solar Activity and Space Weather”.
J.A. Gonzalez-Esparza thanks the support of the CONACyT 48494 and DGAPA-PAPIIT (grant IN1053103) projects. E. Aguilar-Rodriguez thanks the DGAPA-PAPIIT project (grant IN110309-2) and CONACyT
project (grant 101625).
The authors would like to thank the directors and staff of EISCAT and MERLIN for their help and for
the EISCAT and MERLIN data used in this study. EISCAT is supported by the scientific research councils
of China, Finland, France, Germany, Japan, Norway, Sweden, and the United Kingdom (UK). MERLIN is
operated by the University of Manchester (UK) on behalf of the STFC. We would also like to extend particular
thanks to W.A. Coles and B.J. Rickett of UCSD, for making available their analysis routines, expertise, and
advice to us when analysing the EISCAT and MERLIN IPS data. The PHOENIX-2 data were made available
by the Radio Astronomy and Plasma Physics group of the Swiss Federal Institute of Technology, Zurich
(ETZ). The 40 – 800 MHz radio spectral data are from the Type II Burst Catalogue generated and maintained
by the Astrophysikalisches Institut Potsdam. E.A. Jensen thanks L. Jian from the Institute of Geophysics and
Planetary Physics (IGPP), University of California, Los Angeles (UCLA), for input to the CME fits, and also
T.L. Mulligan of The Aerospace Corporation for useful discussions.
All standard SOHO images used in this article are taken from the Coordinated Data Analysis Workshops
(CDAW) CME catalog http://cdaw.gsfc.nasa.gov/CME_list/. The CDAW CME catalog is generated and maintained at the CDAW Data Center by NASA and The Catholic University of America in cooperation with the
Naval Research Laboratory. SOHO is a project of international cooperation between ESA and NASA. The
authors thank and acknowledge the Wind and ACE instrument groups for making their data freely available
for scientific use. The authors also acknowledge the ESA Cluster Active Archive for provision of Cluster data. Geotail magnetic field and plasma data were provided by T. Nagai and Y. Saito through DARTS
at the Institute of Space and Astronautical Science, JAXA. NORSTAR MSP data are courtesy of the University of Calgary. LANL MPA data were provided by Michelle Thomsen, and the IMAGE-FUV data are
courtesy H. Frey and S. Mende. The Polar Cap Index was provided by the Danish Meteorological Institute.
The 13 May 2005 Coronal Mass Ejection
119
The authors would like to extend special thanks to the referees of this manuscript for their welcome
comments, helpful inputs, and the long length of time it will have taken to referee such a long, detailed, and
complex paper.
Open Access This article is distributed under the terms of the Creative Commons Attribution Noncommercial License which permits any noncommercial use, distribution, and reproduction in any medium, provided
the original author(s) and source are credited.
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Mayank Vahia
Tata Institute of Fundamental Research
P. Janardhan
Physical Research Laboratory
Jack Sarfatti
Cornell University
Harijono Djojodihardjo
Institut Teknologi Bandung
