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Solar Physics
DOI: 10.1007/•••••-•••-•••-••••-•
On flare-CME characteristics from Sun to Earth
combining remote-sensing image data with in-situ
measurements supported by modeling
Manuela Temmer1
·Julia Thalmann1
·
Karin Dissauer1
·Astrid Veronig1
·
Johannes Tschernitz1
·
J¨urgen Hinterreiter1
·Luciano Rodriguez2
c
Springer ••••
Abstract We analyze the well observed flare-CME event from October 1, 2011
(SOL2011-10-01T09:18) covering the complete chain of action – from Sun to
Earth – for a better understanding of the dynamic evolution of the CME and its
embedded magnetic field. We study in detail the solar surface and atmosphere
associated with the flare-CME from SDO and ground-based instruments, and
also track the CME signature off-limb from combined EUV and white-light data
with STEREO. By applying 3D reconstruction techniques (GCS, total mass)
to stereoscopic STEREO-SoHO coronagraph data, we track the temporal and
spatial evolution of the CME in interplanetary space and derive its geometry and
3D-mass. We combine the GCS and Lundquist model results to derive the axial
flux and helicity of the MC from in situ measurements (Wind). This is com-
BM. Temmer
manuela.temmer@uni-graz.at
J.K. Thalmann
julia.thalmann@uni-graz.at
K. Dissauer
karin.dissauer@uni-graz.at
A.M. Veronig
astrid.veronig@uni-graz.at
J. Tschernitz
johannes.tschernitz@edu.uni-graz.at
J. Hinterreiter
juergen.hinterreiter@edu.uni-graz.at
L. Rodriguez
rodriguez@sidc.be
1Institute of Physics, University of Graz, Austria
2SolarTerrestrial Center of Excellence, SIDC, Royal Observatory of Belgium
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 1
arXiv:1703.00694v1 [astro-ph.SR] 2 Mar 2017
Temmer et al.
pared to nonlinear force-free (NLFF) model results as well as to the reconnected
magnetic flux derived from the flare ribbons (flare reconnection flux) and the
magnetic flux encompassed by the associated dimming (dimming flux). We find
that magnetic reconnection processes were already ongoing before the start of
the impulsive flare phase, adding magnetic flux to the flux rope before its final
eruption. The dimming flux increases by more than 25% after the end of the
flare, indicating that magnetic flux is still added to the flux rope after eruption.
Hence, the derived flare reconnection flux is most probably a lower limit for
estimating the magnetic flux within the flux rope. We find that the magnetic
helicity and axial magnetic flux are reduced in interplanetary space by ∼50%
and 75%, respectively, possibly indicating to an erosion process. A mass increase
of 10% for the CME is observed over the distance range from ∼4–20 R. The
temporal evolution of the CME associated core dimming regions supports the
scenario that fast outflows might supply additional mass to the rear part of the
CME.
Keywords: CMEs, Flares; Dynamics, Magnetic fields; Corona, Interplanetary
Space, In-situ Data
1. Introduction
Since the launch of the Solar TErrestrial RElations Observatory (STEREO
Howard et al., 2008), the Sun-Earth distance range is well covered as never
before. Having three eyes viewing the Sun from different vantage points, many
new insights about the initiation and subsequent propagation of coronal mass
ejections (CMEs) in interplanetary space could be gained. Novel methods on 3D
reconstructions of CMEs (Thernisien, Vourlidas, and Howard, 2009; Mierla et al.,
2010), and with that, more detailed studies with respect to CME associated
solar surface phenomena (e.g., flares or large-scale waves) were pursued, that
could largely improve the understanding of CMEs (e.g., Kienreich, Temmer,
and Veronig, 2009; Temmer et al., 2010; Patsourakos and Vourlidas, 2012; Bein
et al., 2012). In situ measurements at 1 AU show signatures that can be related to
CME-associated solar surface signatures as well as direct observations of CMEs
in white-light. In this respect, simultaneous on-disk and off-limb observations
provide an invaluable source of linking remote sensing and in situ signatures.
Due to their impact and potential geoeffectiveness, Earth-directed CMEs are of
special interest. Using the unprecedented multi-viewpoint data sets currently
available, we can enhance our knowledge on CME characteristics and their
behavior in interplanetary space. Results obtained using on-disk imagery will
provide valuable information for periods which are limited to single viewpoint
observations.
The close relation between early CME evolution and relation to solar flares
is well acknowledged (e.g., Zhang and Dere, 2006; Temmer et al., 2008). Often
associated to flare-CME events are dark dimming regions observed as decreased
emission in extreme-ultraviolet (EUV) and soft X-rays (SXR). These are most
probably caused by the expansion and evacuation of plasma due to a CME,
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 2
Remote sensing and in-situ CME characteristics
and are therefore interpreted as low-coronal footprints of CMEs (Hudson and
Cliver, 2001). The analysis of dimming regions is of special interest, as the
plasma which is depleted from the corona may reflect the mass which is fed
into the CME, maybe over hours (e.g., see also Zarro et al., 1999; Harra and
Sterling, 2001). Hence, characteristic CME properties may be derived from the
dimming evolution (e.g., Cheng and Qiu, 2016). Qiu et al. (2007) derived the
total magnetic reconnection flux in the low corona for flare-associated CMEs and
their dimming regions, and compared it to the corresponding magnetic flux in
magnetic clouds (MC) probed at 1 AU. For a sample of nine events they found
that the reconnection flux in the flare is related to the magnetic flux of the MC.
However, a straightforward comparison between flare characteristics, or dimming
regions on the solar surface with off-limb measurements or in situ counterparts
is not an easy task as unknown processes causing the dimming (Mandrini et al.,
2007) or projection effects from single spacecraft views (Dissauer et al., 2016)
may lead to erroneous interpretations.
Early studies linking filament and MC characteristics were successfully per-
formed by Bothmer and Schwenn (1998) who related interplanetary magnetic
properties of MCs to filament orientation and handedness at the Sun. The
helicity of an erupting flux rope is assumed to be conserved during the CME
propagation in interplanetary space enabling us to link MCs observed in situ
to their solar sources (e.g. Dasso et al., 2005; Rodriguez et al., 2008). With the
power of multi-spacecraft data, revealing remote sensing as well as in situ data
from different vantage points, we are able to have an even more detailed look
on the different aspects of CMEs, their interplanetary propagation behavior and
associated in situ signatures (e.g., Rodriguez et al., 2011; Kilpua et al., 2013;
M¨ostl et al., 2014). More comprehensively, a variety of case-studies linked in
more detail the different aspects of the Sun to Earth flare-CME events. E.g.
M¨ostl et al. (2008) focused on the comparison between magnetic flux derived
from flare reconnection and in situ data. Bisi et al. (2010) performed an extensive
study using multi-instrument data for the analysis of the CME-associated source
region which was simulated from vector magnetic field data driven by artificial
horizontal flux emergence. In a recent study, Patsourakos et al. (2016) tracked
the cause of a strong space weather event, in particular focusing on the near-Sun
magnetic field strength from which the geoeffectiveness might be assessed.
In the current study we investigate the centrally on-disk located CME-flare
event from October 1, 2011 starting at 09:18 UT (SOL2011-10-01T09:18). Com-
pared to already existing studies, we bring new aspects into the dynamic evo-
lution of a CME and its embedded magnetic field, by analyzing in detail the
solar source region using nonlinear force-free and finite-volume helicity model-
ing, deriving the reconnected flux from the CME associated flare ribbons and
dimming areas. In a novel approach we attempt to combine model results from
3D reconstructions of the CME close to the Sun with in situ models for obtaining
the magnetic field characteristics of the associated MC. We compare the results
derived from remote-sensing imagery and in-situ measurements and discuss the
relationship between the parameters.
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 3
Temmer et al.
2. Data and Methods
We investigate in detail the flare-CME event from October 1, 2011. The CME
event is launched from NOAA active region (AR) 11305 located at N10W08,
associated with a M1.2 GOES class flare (start: 09:18 UT, minor peak 09:37 UT,
major peak: 10:00 UT, end: 10:17 UT).
2.1. Flare energetics
For the flare evolution, we study full disk Hαfiltergrams from Kanzelh¨ohe Obser-
vatory for Solar and Environmental Research (KSO) with a temporal cadence of
about 6 sec covering the time range 09:18 UT until 11:00 UT (P¨otzi et al., 2015).
Together with the information of the magnetic field from the 720s LOS magne-
togram of the Helioseismic and Magnetic Imager aboard the Solar Dynamics
Observatory (SDO/HMI Scherrer et al., 2012; Schou et al., 2012; Hoeksema
et al., 2014), we derive magnetic reconnection rates from the separation of
flare ribbons observed in Hα. The Hαimages are normalized and co-aligned
to the first image (north-up and derotated to the reference time 09:18 UT). The
magnetic field maps are binned to the pixel scale of the Hαfiltergrams using IDL
(coreg map.pro). For the alignment between Hαimages and the magnetograms,
HMI continuum images are used.
As shown in Fig. 1, we derive the flare ribbon separation speed from intensity
profiles calculated along rectangular slices oriented perpendicularly to the PIL
along two directions within each magnetic polarity (tracking paths: N1/2, S1/2).
At each time step the intensity profile of each slice is fitted with a Gaussian func-
tion leading to a distance-time diagram. The time derivative of the polynomial
fit of the derived distance-time curve is calculated to get the ribbon velocity and,
hence, the local reconnection rate (Temmer et al., 2007).
The evolution of the flare ribbons provide us with important information on
the coronal magnetic reconnection process in solar flare-CME events. Assuming
translational symmetry in the flaring arcade that is built up behind the erupting
CME, the reconnected electric field in the corona, Ec, can be derived from the lo-
cal ribbon flare separation speed vraway from the polarity inversion line together
with the underlying normal component Bnof the photospheric magnetic field at
the flare-ribbon location, as Ec=vrBn(cf. Priest and Forbes, 1986). In case the
flare does not occur too far off the disk center, the normal component Bncan
be well approximated by the LOS field as measured by SDO/HMI. Forbes and
Lin (2000) generalized this relation to three dimensions, showing that the rate at
which magnetic flux is swept by the flare ribbons relates to a global reconnection
rate. Assuming that the change of the photospheric field during the flare is small,
this global reconnection rate can be determined from the observations as
˙ϕ(t) = dϕ
dt ≈∂
∂t ZBn(a)da , (1)
with da the newly brightened flare area at each instant and Bnthe normal com-
ponent of the photospheric magnetic field strength underlying da. This relation
basically reflects the conservation of magnetic flux from the coronal reconnection
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 4
Remote sensing and in-situ CME characteristics
1-OCT-2011 09:38:26 UT
Figure 1. KSO Hαfiltergram showing the flare before it reaches its maximum intensity. The
polarity inversion line (PIL) is shown as white line, different directions (tracking paths: N1/2
for the northern part and S1/2 for the southern part) along which the ribbon main motion is
tracked are shown with yellow rectangles, bright flare pixels cumulated until the time of the
image shown (09:38:26 UT) are shown as blue areas (positive polarity) and red areas (negative
polarity).
site to the lower atmosphere, where the flare ribbons are observed (Forbes and
Lin, 2000).
2.2. Coronal dimming
We distinguish two different types of dimming regions, core or twin dimmings
and secondary or remote dimmings (see e.g., Mandrini et al., 2007). Core dim-
mings are found in the form of stationary (long-lived) regions of strongly reduced
EUV emission, and are closely located to the CME eruption site. Being located
in regions of opposite magnetic polarity, they presumably resemble the cross-
sectional area of the erupting flux rope footpoints at low-coronal heights. Remote
dimmings are observed over larger areas, extending to significant distances away
from the eruption site, in the form of reduced EUV emission (though not as
pronounced as in core dimming regions).
We calculate the coronal dimming evolution from SDO/AIA (Pesnell, Thomp-
son, and Chamberlin, 2012; Lemen et al., 2012) data in several wavelengths (most
sensitive to quiet coronal temperatures around ≈0.6−2×106K: 171 ˚
A, 193 ˚
A and
211 ˚
A). The time series covers 12 h from the reference time 09:14 UT. We use
high-cadence (12 sec) observations from 09:14 until 11:14 UT and a successively
reduced cadence (1, 5, and 10 min) for the rest of the time series. The dimming
regions are identified by applying a thresholding technique on logarithmically
scaled base ratio images. A pixel is flagged as a dimming pixel if its logarithmic
relative intensity is lower than −0.5 compared to its pre-event value. As an
indication of the core dimming regions, we use the 10% pixels in the dimming
region that revealed the largest absolute change of their intensity below a certain
threshold intensity. Naturally, applying the thresholding technique to the coronal
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 5
Temmer et al.
images covering different temperature regimes (wavelength bands), results in
different tracked extents of the dimming regions. Visual inspection suggests that
the dimming areas are tracked best in 211 ˚
A which is also supported by previous
studies (e.g., Robbrecht and Wang, 2010; Kraaikamp and Verbeeck, 2015).
For the magnetic field information we use the 720s LOS magnetogram of
SDO/HMI at the beginning of the event. All data were prepared using standard
SolarSoft IDL software (aia prep.pro, hmi prep.pro), filtered for constant expo-
sure time, and differentially de-rotated to the reference time. Based on this, we
study the time evolution of the area of the coronal dimming regions and calculate
the magnetic flux involved in the total dimming regions, only considering pixels
with a magnetic field strength |Bi|>10 G, i.e., above the HMI noise level. We
note that values given for the dimming flux are derived from 211 ˚
A image data
using the arithmetic mean over positive and negative polarity.
2.3. Coronal magnetic field modeling
The 3D coronal magnetic field configuration in and around NOAA 11305 was
modeled based on full-disk vector magnetic field observations from SDO/HMI.
The “hmi.B 720s” data series provides the total field, inclination and azimuth on
the entire solar disk. The azimuth is provided with the 180◦-ambiguity already
resolved in strong-field regions (using a minimum-energy method). For weak-field
regions, we apply a so-called random disambiguation method, using the software
tools provided by JSOC1. From the field, inclination and disambiguated azimuth,
we retrieve the image-plane components of the magnetic field vector, i.e., the
LOS and transverse field. In order to account for projection effects, we de-project
the image-plane data to a heliographic coordinate system, i.e., we derive the true
vertical and horizontal field components, following Gary and Hagyard (1990). A
sub-field of these optimized full-disk magnetic field data, covering the flaring
AR as well as its nearest quiet-Sun surrounding, is used as an input to the
nonlinear force-free (NLFF) coronal magnetic field modeling method (for details
see Wiegelmann and Inhester, 2010, and Sect. 2.2.1 of DeRosa et al. 2015).2
Using the 3D NLFF field as an input, we employ the finite-volume helicity
method of Thalmann, Inhester, and Wiegelmann (2011), in order to estimate
the relative helicity of the CME source region.
2.4. CME morphology and kinematics
For deriving the entire kinematical profile of the CME evolution, we study com-
bined EUV and white-light data from different vantage points using the SECCHI
1The Joint Science Operations Center (JSOC) provides supporting documentation and soft-
ware for HMI data. For details on the data series and related azimuth disambiguation procedure
see http://jsoc.stanford.edu/jsocwiki/FullDiskDisamb.
2We list two important controlling parameters proposed in literature (e.g., Wheatland, Stur-
rock, and Roumeliotis, 2000; Schrijver et al., 2006), in order to quantify the goodness of
the obtained NLFF coronal magnetic field solution. For the current-weighted average of the
sine of the angle between the modeled magnetic field and electric current density, we find
CW sin ≈0.1. For the volume-averaged fractional flux we find h|fi|i ≈ 10−4. (For a perfectly
force-free and solenoidal solution, one would obtain CW sin = 0 and h|fi|i = 0.)
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 6
Remote sensing and in-situ CME characteristics
(Sun Earth Connection Coronal and Heliospheric Investigation) instrument suite
aboard the Solar TErrestrial RElations Observatory (STEREO; Howard et al.,
2008) as well as LASCO coronagraph data aboard SoHO (Brueckner et al., 1995).
On October 1, 2011, the separation angle for STEREO-A–Earth and STEREO-
B–Earth was 104.3◦and 97.5◦, respectively, perfectly suited to derive reliable
CME kinematics for an event centrally located on the solar disk from Earth
view.
For estimating the CME geometry, its main propagation direction, and depro-
jected bulk speed, we use the graduated cylindrical shell (GCS) reconstruction
method (Thernisien, Howard, and Vourlidas, 2006; Thernisien, Vourlidas, and
Howard, 2009), which assumes an ideal flux rope to forward fit the appearance
of the CME in coronagraph white-light data. The CME’s early evolution is
determined by manually tracking the frontal part of the CME along its main
propagation direction, using SECCHI EUVI (with a field-of-view - FoV - up to
1.7 R), COR1 (FoV of 1.4–4.0 R) and COR2 (FoV of 2.5–15.0 R) image
data from STEREO-A and -B. In addition we use the stereoscopic data in order
to calculate the 3D-mass of the CME using the method described in Colaninno
and Vourlidas (2009). Following Bein et al. (2013) we derive the 3D CME mass
evolution corrected for occulter effects over the distance range 1–20 R.
The interplanetary CME propagation from Sun to Earth orbit is tracked
along the main propagation direction of the CME, applying the SATPLOT
software tool3available in IDL SolarSoft. The SATPLOT tool delivers j-maps
for combined COR2, HI1 (FoV 4.0–24.0◦) and HI2 (FoV 18.7–88.7◦) white-light
data making it easy to measure the elongation angle of the CME under study.
The measured elongation-time profile is converted into a radial distance profile
using the propagation direction and angular width, obtained from the GCS
reconstruction. To obtain a range of possible propagation directions, we use
several different conversion methods, including Fixed-Phi (FP), Harmonic Mean
(HM) and Self-Similar Expansion (SSE), as described in Sheeley et al. (1999);
Lugaz, Vourlidas, and Roussev (2009); Davies et al. (2012), respectively. For
calculating the CME speed and acceleration profile from the time-distance data,
we apply the regularization method as described in Temmer et al. (2010).
2.5. In situ CME characteristics
To correctly identify the in situ signatures of the CME (ICME) at Earth orbit,
we perform drag-based-modeling (DBM) in order to simulate its interplanetary
propagation along the main propagation direction (Vrˇsnak and ˇ
Zic, 2007; Vrˇsnak
et al., 2013). As an input, we use the CME’s initial speed, distance, and angular
width, as obtained from the GCS reconstruction. From the results we estimate
the time range most suitable for studying the related ICME characteristics. We
investigate the in situ plasma and magnetic field by using 1-min resolution Wind
data (Lin et al., 1995; Lepping et al., 1995). We apply a Lundquist force-free
cylindrical fit (hereafter referred to as Lundquist model; see Lundquist, 1950)
3http://hesperia.gsfc.nasa.gov/ssw/stereo/secchi/idl/jpl/satplot/SATPLOT User Guide.pdf
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 7
Temmer et al.
Figure 2. (a) Central filament channel as observed in SDO/AIA 304 ˚
A, around the main
sunspot of NOAA 11305, prior to the flare-CME on October 1, 2011. (b) NLFF model magnetic
field lines, outlining the observed filament channel (colors are for better visibility only). The
color coded background resembles the SDO/HMI vertical magnetic field, scaled to ±2 kG. (c)
Orientation of the coronal magnetic field (orange arrows) in a vertical cut through the model
volume, above the path outlined as white solid line in (b). The light and dark violet model
field lines are shown as in (b). (d) Orientation of the coronal magnetic field as in (c), but with
the magnitude of the total electric current density shown as color-coded background.
to the in situ magnetic field data in order to reconstruct the properties of the
ICME’s flux rope, including its interplanetary orientation, radius and axial field
strength. These are then used to calculate its axial magnetic flux and helicity,
following DeVore (2000).
3. Results
3.1. Source region characteristics: pre-flare structure
The initial conditions for the eruption are derived from the pre-flare NLFF
model (employed at 07:59 UT). The NLFF coronal magnetic model shows highly
twisted magnetic fields, along the main PIL (Fig. 2b), that clearly outline the
dark filament observed in AIA 304 ˚
A (Fig. 2a). Projected into a vertical plane,
roughly perpendicular to the main axis of the filament, the coronal magnetic field
vector exhibits a counter-clockwise pattern, i.e., a left-handed sense (orange
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 8
Remote sensing and in-situ CME characteristics
KSO $H!" 09:24$UT$$$ $$$$$$$$$$$$$$$$$$$$$$$$$$$$$$$$ $ 09:59$UT $$$$$$$$$$$$$$$$$$$$$$$$$$$$$$$$$$$$ 10:18$ UT$
a)
d)
b)
c)
e)
f)
Figure 3. Sequences of KSO Hα(top panels) and AIA 304 ˚
A (bottom panels) images dis-
playing the evolutionary stages of the flare, including the impulsive (left), peak (middle) and
decay (right) phase.
arrows in Fig. 2b,c). The NLFF model field lines (violet and purple lines in
Fig. 2c) warp around a central axis, at an approximate height of 3 arc-second
(≈2 Mm) above photospheric level, characterized by strongest electrical current
density (Fig. 2d). These properties are consistent with that of a coronal flux rope
(e.g., Filippov et al., 2015). The total unsigned axial flux within the flux rope
is 1.1×1021 Mx (estimated from the magnetic flux penetrating the vertical
plane shown in Fig. 2c and 2d). Using the 3D NLFF field as an input, we
estimate the relative helicity of the AR core, hosting the flux rope, as HV≈
−3.9×1042 Mx2. The relative helicity is a measure of how much a field is
twisted and/or entangled, with respect to a reference potential field (of vanishing
electric current and helicity). Its sign arises from the negative contribution of the
left-handed fields to the AR’s helicity budget (e.g., review by D´emoulin, 2007).
3.2. Source region characteristics: eruptive phase
3.2.1. Morphology of two-phase filament eruption
Fig. 3 shows the morphology of the flare as observed in Hα(Fig. 3a–c) and
AIA 304 ˚
A (Fig. 3d–f). First, the filament to the south-west of the sunspot is
activated and starts to rise around 09:22 UT (Fig. 3d), simultaneous to an initial
rise of the observed SXR emission to C6 level (compare Fig. 6g). In a second step,
the structures in the south-east of the sunspot are destabilized (around 09:37 UT)
and ending in the final eruption observed around ∼09:59 UT, co-temporal with
the SXR emission rising towards the final M1.2-level. The flare ribbons observed
in Hαshow a consistent evolution. Ribbon formation is observed first to the
west of the sunspot (Fig. 3a), close to the location where the first post-flare
loops become visible in EUV (compare 3d), and evolves towards the south-east
as the flare progresses (Fig. 3b).
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 9
Temmer et al.
09:12-09:24 UT
09:24-09:59 UT
09:59-10:18 UT
a)
c)
e)
b)
d)
f)
Figure 4. Evolution of flare ribbons and coronal dimming, during three different time intervals
covering the early impulsive phase (top), late impulsive phase (middle) and the decay phase
(bottom) of the flare. Right panels: Area covered by core (green filled contours) and remote
(red contours) dimming. The blue rectangle outlines the flare region, shown in the left panels.
Left panels: Locations attributed to flare ribbons (cyan/yellow contours for signatures above
negative/positive photospheric polarity) and core dimming (blue/red filled contours above
negative/positive polarity). The grayscale background resembles the HMI LOS magnetic field
at 09:12 UT, scaled to ±1 kG (left panels) and to ±0.1 kG (right panels) with black/white
color representing the negative/positive polarity, respectively.
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 10
Remote sensing and in-situ CME characteristics
a)
b)
Figure 5. (a) NLFF model field lines calculated from the core dimming regions, tracked
between 09:12 and 09:24 UT (see Fig. 4). The field lines are color coded according their apex
height. The gray-scale background resembles the pre-flare vertical magnetic field, scaled to
±2 kG. (b) NLFF model field lines calculated from the flare pixels tracked in Hαimages for
the same time interval as the core dimming shown in (b). Color-coding of the NLFF field
lines and background as in (b). Only field lines which close within the field-of-view and which
connect photospheric regions of Bz>10 G are shown.
3.2.2. Time evolution of flare-induced ribbons and CME-induced dimming
In Fig. 4, we trace both, the bright flare-induced ribbon emission and the di-
minished emission from the CME-induced dimming. From top to bottom, the
signatures characteristic for the early impulsive (09:12 – 09:24 UT), late impul-
sive (09:24 – 09:59 UT), and decay (09:59 – 10:18 UT) phase of the associated
flare are shown. The panels of the right column outline the time evolution of
the CME-associated coronal dimming, divided into core dimming and remote
dimming (see Sect. 2.2). In the panels of the left column, a close-up of the flare
region (marked by the blue rectangle in the right column) is shown. Here, only the
core dimming (blue and red filled contours) is shown, together with the locations
populated by flare ribbon emission (cyan and yellow contours). During the early
impulsive phase, flare ribbons and core dimming regions (Fig. 4a) appear to the
south-west of the sunspot, coinciding with the location of the filament observed
before the flare (compare Fig. 2a), and marking the footprint of the coronal
magnetic field involved in the first phase of the flare. The flare ribbons and core
dimming evolve towards the south-east of the sunspot only during the second
phase of the flare (the late impulsive phase), coincident with the final eruption
of the filament. With the launch of the CME, the formation of pronounced and
extended remote dimming areas is initiated (see Fig. 4f and compare Fig. 6a–c).
As can be seen, this event reveals a complex interplay between flare brightened
areas and core dimming regions. For this reason we use the core dimming areas
only for qualitative purposes (cf. Sect. 3.3.1).
In Fig. 5a we show NLFF model field lines traced from the detected core dim-
ming area (cf. Section 2.2) during the early impulsive phase (09:12 – 09:24 UT).
This allows us to infer some geometrical properties of the magnetic structure
that later developed into the observed CME. It involves twisted fields (a flux
rope; compare Fig. 2) with apex heights of .7 Mm. A comparison with Fig. 5b)
shows the model field lines traced from the flare pixels tracked within the same
time interval. Besides the low-lying magnetic flux rope to the south-west of the
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 11
Temmer et al.
sunspot, also higher-reaching fields to its south-east (with apex heights up to
≈25 Mm) were subject to magnetic reconnection, demonstrating the magnetic
connection between the different portions of the AR that were involved in the
eruption. Therefore, we can assume that the flare-CME process, initiated in the
form of a filament eruption to the west of the sunspot (coincident with the early
impulsive phase), progressed to the south-east of the sunspot by destabilization
of/reconnection with the overlying magnetic configuration in that part of the
active region (marking the late impulsive phase of the flare).
3.2.3. Relative timing of flare- and CME-associated features
Fig. 6 relates different flare- and CME-associated parameters to distinct phases
during the flare. The two phases of the filament eruption, including the erup-
tion of the filament to the west of the sunspot and the destabilization of the
overlying field to the south-east, match well with the different phases observed
in the GOES X-ray flux and its derivative (black and green curve in Fig. 6g,
respectively), the latter being a proxy for the flare-associated hard X-ray emis-
sion (Neupert, 1968; Veronig et al., 2002). The time derivative of the magnetic
flux associated with the flare ribbon pixels (red and blue curve in Fig. 6d,
for positive- and negative-polarity associated pixels, respectively) reveals major
changes throughout the impulsive phase. On the other hand, the time derivative
of the magnetic flux associated with the dimming pixels (red curve in Fig. 6c)
suggests major changes for the time range covering the CME initiation until the
CME attains maximum speed. At the end of the decay phase, when the SXR
flux decreased again to C-level around 10:45 UT, we find a total accumulated
reconnected flux from the flare ribbon evolution of 2.1·1021 Mx, approximately
two times the flux involved in the dimming (≈1.1·1021 Mx at 10:45 UT).
The local reconnection rate (Fig. 6e and f; deduced from the flare ribbon
separation velocity and associated magnetic flux; cf. Sect. 2.1) is distributed in a
non-uniform way along the flare ribbon (compare the resulting curves along the
two tracking paths N1 and N2 for the northern flare ribbon and along S1 and S2
for the southern flare ribbon, and see also Fig. 1) which has been observed also
in earlier studies (see e.g., Temmer et al., 2007). The (velocity) acceleration time
profile of the CME, as derived from combined EUV and COR1 measurements
(cf. Section 2.4), reveals a close relation with the time evolution of (the derivative
of) the GOES X-ray flux. The reconnected flux associated with the flaring peaks
first, followed by that associated with the dimming, followed by the CME’s
acceleration to its maximum speed (∼550 km s−1at a distance of 4 R; see
Fig. 6b).
3.3. CME 3D characteristics and kinematical evolution
3.3.1. CME 3D-mass and near-Sun kinematics
Fig. 7a presents the near-Sun CME’s 3D-mass evolution corrected for occulter
effects. The CME kinematics up to a distance of ∼20 Ris given in Fig. 7b
(cf. Fig. 6a for the kinematical profile up to ∼5 R). The results suggest that
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 12
Remote sensing and in-situ CME characteristics
Total dimming
Filament
1st - 2nd
S
T
A
R
T
M
A
X
E
N
D
a)
b)
c)
d)
e)
f)
g)
Figure 6. Time evolution of flare-/CME-related parameters, in comparison to the deduced
CME kinematics. a) Distance-, b) velocity- and acceleration-time profile of the CME. c) Re-
connected magnetic flux deduced from dimming pixels (211 ˚
A) and d) from flare ribbon pixels
covering areas of positive (blue) and negative (red) magnetic polarity. Local reconnection rate
deduced along different directions of ribbon motion, normal to the local polarity inversion
line, from the e) Northern and f) Southern flare ribbon. g) GOES soft X-ray flux (black) and
its time derivative (green). Yellow vertical dashed lines mark distinct times within the flare
process, including the nominal flare start, peak (“MAX”), and end time. Gray vertical dashed
lines mark different phases during the filament eruption.
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 13
Temmer et al.
a)
b)
c)
d)
Figure 7. Top to bottom panels : Mass and kinematics of the CME within 20 Rabove the
solar surface. (a) 3D-mass estimate based on COR1 and COR2 observations. A fit (green
dashed line) has been applied to the combined COR1 and COR2 measurements, where only
those masses estimated from COR2 have been used which exceed the COR1-based mass
estimate at 4 R). Based on the fit, the 3D-mass evolution as function of height corrected
for occulter effects (red solid line) and the seed-mass (horizontal black dashed line) is calcu-
lated. (b) CME distance-time evolution as derived from STEREO-A/B white-light images (red
and blue squares, respectively) and the deprojected height of the CME front where 3D-mass
measurements were made (green/red crosses). Co-temporal variation of the (c) core and (d)
total dimming area measured from AIA 171, 193, and 211 ˚
A image date.
the flare-associated ejection had a seed-mass of m0=4.4·1015 g, that increased
with a rate of ∆m=6.1·1013 g R−1
. As a result, we estimate the final mass at a
distance of 20 Ras mend ∼5.5·1015 g.
Assuming that the observed increase in mass is due to a continuous mass flow
that stems from the coronal regions where the CME footpoints are rooted, we
compare the mass evolution to the evolution of the area covered by the CME-
related coronal dimming (Fig. 7c,d). The time evolution of the total dimming
area in three wavelengths (171, 193 and 211 ˚
A; see Fig. 7d), shows an effective
growth starting at ∼09:45 UT that ceased around 10:20 UT. Relatively large
variations in the core dimming area can be traced until ∼11:30 UT.
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 14
Remote sensing and in-situ CME characteristics
Figure 8. GCS modeling results using the simultaneous view from three spacecraft (STERE-
O-B left, LASCO middle, STEREO-A right ) on October 1, 2011. Yellow arrows mark the
CME. The CME is directed to south-east, with a clear tilt with respect to the ecliptic plane.
The FoV of LASCO C2 covers 2–6 R, STEREO-A/B 2.5–15 R.
3.3.2. CME geometry
Table 1. CME characteristics near the Sun and near Earth, using GCS modeling for the CME
geometry, 3D-mass calculation, and DBM to derive the CME’s propagation characteristics in
interplanetary space and arrival at Earth.
Near Sun Near Earth (1 AU)
GCS source region E08S08 GCS apex radius 1.7·1012 cm (0.11 AU)
GCS tilt vs. ecliptic 45◦GCS flux rope length (L) 3.9·1013 cm (2.6AU)
GCS face-on width 72◦GCS volume 9.0×1037 cm3
GCS edge-on width 34◦Density* 25–35 cm−3
GCS 3D speed 450 km s−1DBM arrival time 2011-10-05 07:37 UT ±5 h
3D-mass 5.5·1015 g DBM impact speed 426 km s−1±30 km s−1
*The density is calculated under the assumption that the mass stays constant beyond
20 Rand is uniformly distributed within the derived CME volume.
Having three different vantage points, with almost perpendicular separation
angles of the two STEREO satellites with respect to SoHO in the Sun-Earth
line, we are able to reconstruct the 3D geometry of the CME from white-light
coronagraphic data. The lower panels of Fig. 8 show the best fit (green cones),
resulting from the GCS 3D flux rope model, when requiring that the boundary of
the GCS model flux rope match the outer edge of the CME shape (indicated by
yellow arrows in the upper panels) in STEREO-B (left), LASCO (middle) and
STEREO-A (right) white-light images. At t0=13:30 UT we obtain from the GCS
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 15
Temmer et al.
model a CME distance of r0=12 R, a propagation direction of φCME =−5◦, a
speed of v0=450 km s−1, and a CME half-width of λ=26◦. Note that due to
the tilt of the reconstructed CME body (∼45◦), we take an average of the face-
and edge-on half-width. All parameters derived from the GCS modeling are
summarized in Table 1.
The obtained values are used as input for the DBM, in order to model the
CME’s interplanetary propagation. This allows us to compare distinct model
parameters with actual in situ measurements at Earth orbit. Using a drag value
of γ=0.2×10−7and an ambient solar wind speed of w=380 km s−1, we estimate
the ICME to arrive at Earth on 2011 October 5 at 07:37 UT (±5 h), with an
impact speed of 426 km s−1(±30 km s−1). Comparison with Wind observations
allows us to determine the arrival of the CME-associated shock at 07:36 UT,
with an impact speed of ∼460 km s−1and followed by a magnetic structure
lasting from ∼10–22 UT (cf. Fig. 10). The ICME caused a moderate geomagnetic
storm of Dst=−43 nT (Richardson and Cane, 2010, “R&C List”4). The modeled
and measured results show a quite good match revealing that the CME only
marginally decelerated on its way from Sun to Earth.
3.3.3. Full kinematical profile
We were able to deduce the full kinematical profile of the CME all the way from
the low solar corona up to 1 AU, based on combined EUV and white-light data.
Fig. 9 shows the track of the CME in interplanetary space (covered by COR2,
HI1, and HI2 data). By applying well-established fitting routines and assuming a
constant propagation speed, we deduce the direction of propagation as E15±10
(SSE:−15◦, HM:−7◦, FP:−22◦; see top right panel in Fig. 9). Importantly, this
result is in accordance with the direction of propagation derived from GCS
modeling, so that we can safely use the value of E15 for the conversion of the
measured elongation angle to radial distances, and thus, for deriving the (I)CME
kinematics, including the speed and acceleration profiles (see Fig. 10a–c and
Section 2.4 for details). The CME front as observed in HI1+2 cannot be entirely
tracked to the distance of L1, however, inspecting Fig. 10a, we see that a linear
extrapolation of the derived kinematics would match well with the arrival of the
CME at Wind spacecraft.
The plasma and magnetic field properties measured in-situ by Wind, covering
the time range 2011 October 4 00:00 UT to October 6 24:00 UT, are shown
in Fig. 10d–i. They suggest that the CME shock-sheath structure arrived at
Earth on October 5 at 07:36 UT (indicated by the blue dashed vertical line).
Signatures typical for a MC (Burlaga, 1991) were observed, including (i) a
rotating magnetic field vector (between 10:00 UT and 22:00 UT; see Fig. 10e–
g), (ii) an enhanced magnetic field strength (Fig. 10d), and (iii) a temperature
below the typical quiet solar wind temperature (Richardson and Cane, 1995).
Applying a Lundquist model to the in situ measured data, we deduce an axial
field strength of B0=12.1 nT, a radius of the MC of r0=1.75·1012 cm, and a
4http://www.srl.caltech.edu/ACE/ASC/DATA/level3/icmetable2.htm
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 16
Remote sensing and in-situ CME characteristics
Figure 9. Left: Interplanetary propagation of the CME under study (red line) tracked using
SATPLOT j-maps. Top right: Conversion results from the derived elongation angle using
several methods with different assumptions on the CME geometry (FP, HM, SSE - for more
details see Sect. 2.4). Bottom right: DBM graphical output (swe.uni- graz.at) using as initial
values the parameters derived from GCS model fit.
relative orientation of the MC in interplanetary space (the axis of the embedded
flux rope being inclined ≈60◦with respect to the Sun-Earth line and ≈54◦with
respect to the ecliptic plane). Importantly, the inclination with respect to the
ecliptic as well as the estimated radius of the MC agree with the corresponding
value obtained from GCS modeling (cf. Table 1 and Section 3.3.2).
In an original approach, we combine the results obtained from the GCS mod-
eling at 1 AU (cf. Table 1) and the cloud parameters derived from Lundquist
model of the in situ data (cf. Table 2), in order to compute the axial flux, Φax ,
and helicity, H, of the MC following DeVore (2000):
Φax = 1.4·B0·r2
0,(2)
and
H= 0.6·Hs·B2
0·r3
0·L. (3)
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 17
Temmer et al.
Remote data
In situ data
(Wind)
a)
b)
c)
d)
e)
f)
g)
h)
i)
Figure 10. Top three panels: CME kinematics covering the distance range from Sun to the
arrival of the CME at Wind spacecraft (blue dashed vertical line), including the (a) measured
CME front height from the solar surface (black plus signs), (b) velocity, and (c) acceleration
as a function of time. Bottom six panels: In situ measurements from Wind for the CME
showing the magnitude (d) and direction (e–g) of the local interplanetary magnetic field in
GSE coordinates. The results of the Lundquist fit applied to the in-situ measurements of the
MC covering the time span October 5, 2011 10:00–22:00 UT are indicated by red solid lines.
(h) Proton speed (solar wind bulk speed; black line), proton density (red line), and (i) proton
temperature (black line), together with the expected temperature Texp for quiet solar wind
conditions (red line), based on which the extension of the MC has been determined (brown
dashed vertical lines).
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 18
Remote sensing and in-situ CME characteristics
Table 2. In situ characteristics using a cylindrical force free model fit (see Fig. 10).
For comparison we list the parameter values as derived from NLFF model results and
the solar source region studies.
In situ measurements (Wind)
Shock arrival time 2011-10-05 07:36 UT R&C list
Impact speed 460 km s−1R&C list
MC start 2011-10-05 10:00 UT R&C list
MC end 2011-10-05 22:00 UT R&C list
Lundquist model results*
Axial field magnitude 12.1 nT B0
MC radius (GSE) 1.75·1012 cm (0.12 AU) r0
φ60.7◦angle to Sun-Earth line
Θ 54.1◦angle with ecliptic
In situ CME magnetic characteristics
Φax 5.2·1020 Mx
H−1.8·1042 Mx2
Solar flare-CME magnetic characteristics
Φax 1.1×1021 Mx NLFF
HV−3.9×1042 Mx2NLFF
Ribbon flux (accumulated) 2.1·1021 Mx Hα(09:18–10:45 UT)
Dimming flux (accumulated) 1.1·1021 Mx 193 ˚
A(09:18–10:45 UT)
Dimming flux (accumulated) 1.4·1021 Mx 193 ˚
A(09:18–11:30 UT)
(*)We apply the Lundquist model to all in situ data between 09:50 and 22:00 UT on
October 05, 2011.
Here, Hsdenotes the helicity sign, which is set to −1, corresponding to the
left-handed flux rope deduced from the in situ observed ICME signature (see
Fig. 10d–i) and Lis the length of the MC that is calculated from the circum-
ference of the GCS model result viewed face-on at 1 AU (see Table 1). As a
result, we obtain for the MC Φax=5.2·1020 Mx and H=−1.8·1042 Mx2, in basic
agreement within a factor of two with the corresponding values derived for its
source region on the Sun (summarized in Table 2 and see also Section 3.1).
4. Discussion and conclusion
We study in detail the CME event from October 1, 2011. The analysis includes a
wealth of data combining remote sensing and in situ instruments to investigate
the complete chain of action for the CME eruption and its evolution from Sun
to Earth. We obtain detailed information on the solar surface signatures of the
associated flare, magnetic field characteristics, and dimming regions that are
subsequently related to the in situ plasma and magnetic field properties of the
CME.
The flare-CME event is associated with a filament eruption that actually
occurred in two-steps, starting in the west of the source AR 11305 and moving
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 19
Temmer et al.
towards south-east. The NLFF results can well explain the process and demon-
strate the magnetic connection, showing that besides the low-lying magnetic flux
rope to the south-west of the sunspot, also higher-reaching fields to its south-east
were subject to magnetic reconnection. This is reflected in the CME propagation
direction (E15) when compared to the source region coordinates (W08), and also
in the location of the remote dimming regions. We derive that the magnetic flux
rope of the CME is fed by two components, low-lying twisted magnetic fields
(rooted in core dimming regions) and sheared overlying magnetic fields (rooted
in flare pixels) involved in the eruption. We derive a flare reconnection flux of
2.1·1021 Mx and a dimming flux of 1.1·1021 Mx.
From Hαemission we obtain the magnetic flux injected to the CME flux rope
at different stages of eruption. We derive an equal amount of flare reconnection
flux during the first impulsive phase of the flare (09:18–09:44 UT), i.e. before
the SXR emission reaches M-level, and during the flare’s major impulsive phase
(>09:44 UT). Hence, reconnection processes were well ongoing before the fila-
ment started to erupt (09:37 UT) followed by the restructuring of the magnetic
field. In comparison, the dimming flux, covering remote areas in the outskirts
of the AR, shows regions involved in the reconnection process at a later time
when the CME has already fully erupted (cf. bottom panel of Fig. 4). Over the
time range 10:45–11:30 UT, hence, after the flare has ceased, the dimming flux
increased from 1.1·1021 Mx to 1.4·1021 Mx. This indicates that magnetic flux
might have been added to the flux rope due to ongoing magnetic restructuring,
too weak to produce visible Hαflare ribbon emission. Therefore, the value
of 2.1·1021 Mx for the total axial flux is most probably a lower limit. When
comparing this to the in situ axial magnetic flux of the MC (∼0.5·1021 Mx)
we find that it is reduced by at least 75% and that the helicity is reduced by
a factor of two. This might refer to an erosion of the MC while propagating
in interplanetary space (e.g., Dasso et al., 2006; Ruffenach et al., 2015). For
calculating the helicity from in situ data, we took the best estimate of the MC
length L, as derived from GCS modeling.
The determination of the magnetic flux in situ as well as the MC radius
rand Lis prone to substantial errors. This is because the parameters used
in its determination are derived from the fitting of an idealized magnetic field
model (Lundquist force free cylindrical fit in our case) to in situ data, where
the selection of the MC boundaries affect the calculations. Furthermore, the
in situ models rely on a single 1D spacecraft crossing through a 3D structure,
taking many assumptions into play (D´emoulin, Janvier, and Dasso, 2016). Also
the GCS model is a fit of an idealized shaped CME to white-light data. In this
respect we note that the MC radius, an important parameter to calculate the flux
and helicity respectively, as determined from the in situ model fit matches well
with the radius as derived from the GCS extrapolated to 1 AU. For other case
studies on this issue including poloidal flux components we refer to Mandrini
et al. (2005); Attrill et al. (2006); Qiu et al. (2007).
The temporal profile of core dimming areas gives indication that the ob-
served CME mass increase of 10% in total, is supplied by the fast outflow from
the core dimming regions. The CME mass consists of coronal plasma which
gets compressed and moved away from the Sun due to the explosive release of
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 20
Remote sensing and in-situ CME characteristics
magnetic energy. The associated dimming regions map the evacuation of the
plasma, with the core dimming marking the CME footpoints and the remote
dimming the CME body. In a qualitative approach we attempt to relate the
temporal evolution of the core dimming and the CME mass increase. The plasma
evacuated from the core dimming area would be detected in COR2 white-light
only beyond the occulter radius of 2.5 R. Assuming an outflow speed of the
order of 100–200 km s−1(e.g., Zarro et al., 1999; Harra and Sterling, 2001; Tian
et al., 2012), the plasma flow could be a detectable part of the CME mass after
∼1.5–3 h (assuming a detection height beyond 4 R, see Bein et al., 2013, this
would yield 3–5.75 h). EUV observations reveal that the major changes in the
core dimming ends around 11:30 UT, accordingly, outflows would feed mass into
the CME until that time. Taking the propagation time as described above into
account, this mass would become visible in the coronagraph white-light data
latest around ∼17:15 UT. The CME apex is at a distance of about 17–18 Rat
that time. This is consistent with statistical results showing that the increase in
CME mass is primarily supplied to the rear-part of the CME to distances below
20 R(see Bein et al., 2013). This is supported by Bemporad and Mancuso
(2010) who observed at a distance of 4.1 Rcontinuous outflows in UVCS data
over hours after the CME shock propagated through. They concluded that the
transit of the CME flank left the coronal magnetic field open over ∼6 hours,
facilitating fast plasma outflow before the corona recovered to the pre-CME
configuration slowing down the outflowing plasma.
We calculate the CME density, using the GCS volume derived for the CME
apex to be at 1 AU, and the observed 3D-mass at 20 R, which is assumed
to be conserved. The derived density (25–35 cm−3) is comparable within a
factor of two to the in situ measurements (10–15 cm−3). However, the unknown
plasma distribution in a CME volume still leaves many questions open, such
as compression during the eruption and subsequent expansion as well as mass
supply from fast outflows like core dimming regions. We note that the CME
mass increase might continue during propagation in interplanetary space due
to material swept up from the solar wind (“snowplow effect”; Cargill, 2004). In
the literature one finds quite high factors of the order of 2–3 for the CME mass
increase in interplanetary space (e.g., Lugaz, Manchester, and Gombosi, 2005;
DeForest, Howard, and McComas, 2013). This might have effects on the drag
that the CME experiences during its propagation phase. According to the DBM
results from the event under study, the drag was of “normal” type and from the
density estimate and comparison to the in situ plasma density data, we could
not confirm a substantial mass increase.
Combining model data at various distance ranges gives us new insight into
the CME characteristics as it propagates from Sun to Earth. However, the
uncertainties especially in the derived magnetic field parameters and the lack
of in situ data at close distances to the Sun still leaves many questions open.
New missions such as Solar Orbiter or Solar Probe Plus will deliver most eligible
data sets to further pursue such studies.
Acknowledgments We thank the anonymous referee for helpful comments. The study was
funded by the Austrian Space Applications Programme of the Austrian Research Promotion
SOLA: temmer_revised2.tex; 15 October 2018; 12:13; p. 21
Temmer et al.
Agency FFG (ASAP-11 4900217) and the Austrian Science Fund (FWF): P24092-N16 and
P25383-N27. J.K.T. acknowledges the excellent collaboration within the International team on
Magnetic Helicity at the International Space Science Institute (ISSI, Bern). L.R. acknowledges
support from the Belgian Federal Science Policy Office through the ESA-PRODEX program.
The presented work has received funding from the European Union Seventh Framework Pro-
gramme (FP7/20072013) under grant agreement No. 606692 [HELCATS]. This research was
partially funded by the Interuniversity Attraction Poles Programme initiated by the Belgian
Science Policy Office (IAP P7/08 CHARM). We thank A. Gulisano and M. Leitner for some
helpful discussions.
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