Off-limb Solar Coronal Wavefronts From SDO/AIA EUV Observations---Implications For Particle Production
K. A. Kozarev\altaffilmark{1
K. E. Korreck
V. V. Lobzin
M. A. Weber
N. A. Schwadron
We derive kinematic properties for two recent solar coronal transient waves observed off the western solar limb with the Atmospheric Imaging Assembly (AIA) onboard the Solar Dynamics Observatory (SDO) mission. The two waves occurred over $\sim10$-min intervals on consecutive days---June 12 and 13, 2010. For the first time, off-limb waves are imaged with a high 12-s cadence, making possible detailed analysis of these transients in the low corona between $\sim1.1$--2.0 solar radii ($R_{s$). We use observations in the 193 and 211~{\AA AIA channels to constrain the kinematics of both waves. We obtain initial velocities for the two fronts of $\sim1287$ and $\sim736$ km~s$^{-1$, and accelerations of $-1170$ and $-800$ m~s$^{-2$, respectively. Additionally, differential emission measure analysis shows the June 13 wave is consistent with a weak shock. EUV wave positions are correlated with positions from simultaneous type II radio burst observations. We find good temporal and height association between the two, suggesting that the waves may be the EUV signatures of coronal shocks. Furthermore, the events are associated with significant increases in proton fluxes at 1~AU, possibly related to how waves propagate through the coronal magnetic field. Characterizing these coronal transients will be key to connecting their properties with energetic particle production close to the Sun.\\\\
ms_v4.tex
%%
%% Beginning of file 'sample.tex'
%%
%% Modified 2005 December 5
%%
%% This is a sample manuscript marked up using the
%% AASTeX v5.x LaTeX 2e macros.

%% The first piece of markup in an AASTeX v5.x document
%% is the \documentclass command. LaTeX will ignore
%% any data that comes before this command.

%% The command below calls the preprint style
%% which will produce a one-column, single-spaced document.
%% Examples of commands for other substyles follow. Use
%% whichever is most appropriate for your purposes.
%%
\documentclass[12pt,preprint]{aastex}
%%\includegraphics
%% manuscript produces a one-column, double-spaced document:

%%\documentclass[manuscript]{aastex}

%% preprint2 produces a double-column, single-spaced document:

%\documentclass[preprint2]{aastex}

\usepackage{natbib}
\usepackage{amsfonts}

\usepackage{color}
\usepackage{ulem}

\bibliographystyle{apj}

%% Sometimes a paper's abstract is too long to fit on the
%% title page in preprint2 mode. When that is the case,
%% use the longabstract style option.

%% \documentclass[preprint2,longabstract]{aastex}

%% If you want to create your own macros, you can do so
%% using \newcommand. Your macros should appear before
%% the \begin{document} command.
%%
%% If you are submitting to a journal that translates manuscripts
%% into SGML, you need to follow certain guidelines when preparing
%% your macros. See the AASTeX v5.x Author Guide
%% for information.

%\newcommand{\vdag}{(v)^\dagger}
%\newcommand{\myemail}{skywalker@galaxy.far.far.away}

%% You can insert a short comment on the title page using the command below.

%\slugcomment{Not to appear in Nonlearned J., 45.}

%% If you wish, you may supply running head information, although
%% this information may be modified by the editorial offices.
%% The left head contains a list of authors,
%% usually a maximum of three (otherwise use et al.).  The right
%% head is a modified title of up to roughly 44 characters.
%% Running heads will not print in the manuscript style.

\shorttitle{Coronal Limb EUV Waves and Particles}
\shortauthors{Kozarev K. A., Korreck K. E., Lobzin V.V., Weber M. A., Schwadron N. A.}

%% This is the end of the preamble.  Indicate the beginning of the
%% paper itself with \begin{document}.

\begin{document}

%% LaTeX will automatically break titles if they run longer than
%% one line. However, you may use \\ to force a line break if
%% you desire.

\title{Off-limb Solar Coronal Wavefronts From SDO/AIA EUV Observations---Implications For Particle Production}

%
%OFF-LIMB CORONAL WAVES FROM HIGH-CADENCE EUV AND RADIO OBSERVATIONS - KINEMATICS AND MORPHOLOGY
%% Use \author, \affil, and the \and command to format
%% author and affiliation information.
%% Note that \email has replaced the old \authoremail command
%% from AASTeX v4.0. You can use \email to mark an email address
%% anywhere in the paper, not just in the front matter.
%% As in the title, use \\ to force line breaks.

\author{K. A. Kozarev\altaffilmark{1}}
\affil{Astronomy Department, Boston University,
    Boston, MA 02215}
\email{kamen@bu.edu}

\author{K. E. Korreck}
\affil{Harvard-Smithsonian Center for Astrophysics,
    Cambridge, MA 02138}
%\sout{60 Garden Street,} 

\author{V. V. Lobzin}
\affil{School of Physics, University of Sydney, 
 NSW 2006, Australia}
% \sout{School of Physics A28,}

\author{M. A. Weber}
\affil{Harvard-Smithsonian Center for Astrophysics,
 Cambridge, MA 02138}
% \sout{60 Garden Street,}
\and

\author{N. A. Schwadron}
\affil{Institute for the Study of Earth, Oceans, and Space, University of New Hampshire, 
 Durham, NH 03824}
% \sout{8 College Road,}


%% Notice that each of these authors has alternate affiliations, which
%% are identified by the \altaffilmark after each name.  Specify alternate
%% affiliation information with \altaffiltext, with one command per each
%% affiliation.
\altaffiltext{1}{also at Harvard-Smithsonian Center for Astrophysics,
Cambridge, MA 02138}
%\sout{60 Garden Street,}

\begin{abstract}
We derive kinematic properties for two recent solar coronal transient waves observed off the western solar limb with the Atmospheric Imaging Assembly (AIA) onboard the Solar Dynamics Observatory (SDO) mission. The two waves occurred over $\sim10$-min intervals on consecutive days---June 12 and 13, 2010. For the first time, off-limb waves are imaged with a high 12-s cadence, making possible detailed analysis of these transients in the low corona between $\sim1.1$--2.0 solar radii ($R_{s}$). We use observations in the 193 and 211~{\AA} AIA channels to constrain the kinematics of both waves. We obtain initial velocities for the two fronts of $\sim1287$ and $\sim736$ km~s$^{-1}$, and accelerations of $-1170$ and $-800$ m~s$^{-2}$, respectively. Additionally, differential emission measure analysis shows the June 13 wave is consistent with a weak shock. EUV wave positions are correlated with positions from simultaneous type II radio burst observations. We find good temporal and height association between the two, suggesting that the waves may be the EUV signatures of coronal shocks. Furthermore, the events are associated with significant increases in proton fluxes at 1~AU, possibly related to how waves propagate through the coronal magnetic field. Characterizing these coronal transients will be key to connecting their properties with energetic particle production close to the Sun.\\\\
\end{abstract}



%IMPULSIVE VS. GRADUAL EVENTS??? WHAT HAPPENS TO PARTICLE SIGNATURES WHEN YOU OBSERVE A WELL-CONNECTED SHOCK IN THE CORONA THAT DECELERATES FAST???
\section{INTRODUCTION}

%WHAT IS THE POINT OF THIS PAPER?
%THE POINT IS TO TALK ABOUT CORONAL OFF-LIMB EUV SHOCK WAVE SIGNATURES IN ULTRA-HIGH CADENCE AIA IMAGING. THE PAPER IS AN INITIAL ATTEMPT TO ANSWER THESE TWO QUESTIONS:
%I. WHAT ARE THE REMOTE MORPHOLOGICAL AND DYNAMIC SIGNATURES OF EUV CORONAL SHOCK WAVES? 
%II. WHAT ARE SOME OBSERVABLE PARAMETERS THAT COULD TELL US WHETHER THE SHOCK IS EFFECTIVE IN PRODUCING ENERGETIC PROTONS?

%1. PRESENT PROPERTIES OF TWO CORONAL (SHOCK) BRIGHT FRONTS IN UV - 
%1.1. KINEMATICS - VELOCITY, ACCELERATION - WHY IS THIS IMPORTANT? 
%1.2. MORPHOLOGY OF THE EMISSION
%-Multiple ripples seen in AIA, radio type II - corresponding to acceleration in 'bays'?
%-Widening of the sheaths between the CME and the front.
% 
%Some questions:
%Why is the post-shock sheath not uniform in intensity (there seems to be brightening along the northern flanks)? What does this imply for the upstream plasma conditions? Higher intensity must mean more plasma heating/compression, so if the shock width and speed is the same, does that mean the Alfven speed is different, or perhaps that there is more pressure upstream?
%Why does the post-shock sheath change in width? Is there a connection between the sheath width, the sheath intensity, and the shock strength? Is this a function of position along the shock front, a function of time?

%1.3. TEMPERATURE RESPONSE OF AIA CHANNELS - DEM from Mark;
%In the absence of those, take two limiting cases:
% - Assume only temperature changes(n=const). Discuss my attempts to reconstruct isothermal EM a la Schmelz et al. (2010)
% - Assume only density changes(T=const). Then the intensity is directly related to density, and the ratio of emissions at same wavelength but at different times gives the square of the density ratio.


%2. RADIO SIGNATURES SIGNIFY THAT THERE ARE SHOCKS, AND THEY ARE EFFICIENT AT ACCELERATING ELECTRONS, AND PERHAPS ALSO PROTONS. THUS, EXAMINING THESE FEATURES THAT ARE ASSOCIATED WITH SHOCKS AT VERY GOOD SPATIAL AND TEMPORAL RESOLUTION.


%INTRODUCTION:
%1. WHY OBSERVE CORONAL WAVES? - THEY MAKE ENERGETIC PARTICLES.
A growing body of theoretical and observational research suggests that charged solar energetic particles (SEPs) gain most of their energy at traveling shocks relatively close to the Sun \citep{Zank:2008}. Interplanetary shocks have been well studied with in-situ measurements near Earth and throughout the solar system \citep{Stone:1985,Forbes:2006}. Many SEP bursts observed close to Earth are not directly associated with Earth-detected shocks. This suggests that SEPs are accelerated much closer to the solar corona, possibly by shocks near the Sun \citep{Reiner:2007}. Coronal shocks could accelerate particles to very high energies in short periods \citep{Roussev:2004}. However, field and shock geometry are key parameters in the ability of shocks to accelerate particles regardless of the shock strength, especially near the Sun \citep{Giacalone:2006}.
%The difficulty with forming shocks close to the Sun is that the medium is strongly mediated by magnetic fields causing high characteristic speeds \citep{Evans:2008}. 
%On the other hand, it seems that not just any shock can accelerate particles, but only stronger ones (\cite{Shen:2007}, \cite{Kocharov:2009}.)



%3. WHY ULTRA-HIGH CADENCE? ASSUMING THAT SHOCKS ARE MOST EFFECTIVE IN ACCELERATING PARTICLES CLOSE TO THE SUN, WE WANT TO KNOW THEIR EARLY STAGES OF DEVELOPMENT, SINCE THEY HAVE MUCH ENERGY TO GIVE AWAY, AND HIGH PLASMA DENSITIES FROM WHICH TO DRAW A SEED POPULATION FOR ACCELERATION. EVEN MORE SIMPLY - MOST SHOCKS ARE FAINT AND FAST, SO WE NEED HIGH CADENCE JUST TO DETECT THEM - IT IS EXTREMELY HARD TO DETERMINE WAVE/SHOCK PARAMETERS JUST BASED ON TWO-THREE POINTS...

%4. HAS THIS BEEN DONE BEFORE?
Coronal shocks have been observed earlier\citep{Pick:2006,Nindos:2008}. \citet{Maia:2000} reported on fast coronal transients propagating with similar speeds in both radio and white light. \citet{Vourlidas:2003} used coronagraph observations to study a white-light coronal shock beyond 2.5 $R_S$. Recent results \citep[e.g.,][]{Gallagher:2010, Patsourakos:2009a, Patsourakos:2009b} have used EUV observations to show the intimate connection between EUV waves and CMEs. However, there is still considerable debate about how shocks appear in these observations. Additionally, a widely used means of characterizing coronal shock kinematics is observations of drifting metric radio emissions from the Sun (approximately 18--180 MHz). These type II radio bursts are associated with coronal shocks accelerating electrons that excite plasma radio emissions \citep{McLean:1985,Reiner:2003,Mancuso:2003}.

%5. OUTLINE OF THE PAPER
Ultra-high cadence EUV imaging observations of off-limb coronal waves are presented, for two recent solar eruptions on June 12 and June 13, 2010. We use the Atmospheric Imaging Assembly (AIA) instrument \citep{Title:2006} aboard the Solar Dynamics Observatory. The temporal resolution, $\sim$12~s, allows for following the evolution of impulsive features in the low corona (1.2--2.0~$R_S$)---a capability newly available in EUV imaging of the Sun. About 3 hours after the June 12 event and 2 hours after the June 13 event, elevated proton fluxes (~6.5 MeV) were observed at 1 AU, leading us to investigate the connection between the remote wave observations and in-situ particle fluxes. We combine simultaneous EUV wave and type II radio burst observations with a coronal magnetic field model to investigate the morphology, kinematics, thermal and density properties of the wavefronts, and their energetic particle production capability.

%There have been numerous studies of CMEs and associated shocks observed in the ultraviolet, with spectrographs \citep{Raymond:2000}, radio \citep{Mancuso:2008,Bemporad:2010}, and with white light coronagraphs \citep{Reiner:2003}. While radio observations have high cadence ($<3$ s), optical observations have cadence on the order of a minute, which is far too slow for characterizing explosive plasma events moving at speeds of hundreds to thousands km~s$^{-1}$ in the low corona. An exception is spectroscopic observations, which have high temporal cadence and are used to determine shock properties accurately from UV emission line ratios. However, the spatial extent of solar spectrographs is very limited, so they cannot monitor the entire Sun or even an extended region, but have to rely on `chance' observations of coronal shocks. Cadence is no longer a limitation after the launch of the Solar Dynamics Observatory (SDO). 

The Letter is structured as follows: In Section~2, we detail the AIA and radio observations used. In Section~3 the kinematics, morphology, and physical properties of the EUV transients are described. We summarize our findings in Section~4.



\section{OBSERVATIONS AND DATA ANALYSIS}

\subsection{EUV observations}

%For the EUV observations we used two bandpass channels of the SDO/AIA telescopes. The first is sensitive mostly to coronal Fe XII ion emission (log[T]=6.1), with peak response at 193~\AA. The second bandpass is sensitive to coronal Fe XIV ions (log(T)=6.3), peaking at 211~\AA.
We used observations from two AIA channels peaking at 193(FeXII) and 211~{\AA}(FeXIV). We refer to them as the 193 and 211 channels throughout the paper. Both channels have $\sim12$~s cadence, $~6$-s lag between the two. The data were processed to level 1.5 using a standard AIA pipeline. Base difference images were produced from an average of ten subframes immediately preceding the events. Event movies can be found as online supplemental materials to this Letter. For temperature and density analysis of the June 13 event, we use six EUV AIA channels dominated by Fe lines (details are given below).
% For both events, we selected a region of 1024x1024~pixels out of the 4096x4096-pixel image, roughly centered on the active region (AR) of interest, and extending to the edge of the AIA field of view (FOV). These subframes were chosen to contain most of the features in both events at all times.
%This ensured that long-lived features (features persisting for more than several minutes) are removed, while enhancing any transients. Base differencing was done for each channel and each event. 

The first event occurred on June 12, 2010 above active region (AR) 11081 located close to the northwestern limb (N23W43). The EUV transient coincided with an M2.0 X-ray flare between 00:30--01:02 UT, peaking at 00:57 UT. We considered observations between 00:56--01:03~UT, the times during which we could detect and measure the coronal transient feature in the FOV of the AIA instrument. During this period a faint, but discernible front was launched roughly radially above the AR.

The second event occurred on June 13, 2010 above AR11079, on the southwestern limb (S25W84). It coincided with an M1.0 flare between 05:30-05:44~UT, peaking at 05:39 UT. In EUV an eruption started at 05:34~UT on the limb, turning into a CME loop propagating radially outward. At 05:37~UT, a hemispheric wavefront appeared (in both 193 and 211~channels) in front of the CME and separated from it, traveling in the same direction but markedly faster. The wave reached the AIA FOV edge at 05:42 UT, followed by the CME at 05:44 UT. %We believe the faster feature was a separate wavefront, not part of the CME, since it appears abruptly as a wide front and cannot be traced to the flare region.

In Figure \ref{fig1}, panels A and C show two base difference images in the 211 AIA channel of the June 12 and 13 events, respectively. Dashed lines trace the wavefronts. Although both events were only clearly visible in difference images, the second event was notably brighter, exhibiting lower velocities, as we show below. 
%Another obvious difference between the two waves is the possible lack of a driver behind the wavefront in the June 12 front (traced in panel C as dotted lines for the June 13 wave.)

Since wave signatures were very faint, we made manual measurements. For each subframe in each event, the expanding wavefront edge was selected along three radial profiles close to the wave's nose. To reduce measurement errors, they were repeated ten times for each image sequence in both channels. We fitted second-order polynomials to measured positions in order to obtain front velocities and accelerations, using MPFIT routines \citep{Markwardt:2009} combined with a statistical bootstrapping technique \citep{Efron:1979}. Since the waves were very dim, we only managed meaningful observations for two profiles in each event. Third-order fits were also attempted, but did not produce significantly different results.

We also corrected for plane-of-sky projection of the wavefronts, assuming spherical waves propagating radially away from the Sun. Then, the brightest EUV emission is detected at front edges. We deprojected front positions by assuming $r=r'/\sin(\phi)$, where $r$ and $r'$ are the true and projected radial distances from the flare site, respectively, and $\phi$ is the AR heliographic longitude. Velocities and accelerations are presented in Table \ref{table1}.

\subsection{RADIO OBSERVATIONS}
Metric radio spectra were provided by the Learmonth Solar Radio Observatory (Western Australia). Type II bursts indicate electron acceleration by coronal shocks, which  may also accelerate protons and heavier ions. The Newkirk coronal electron density model \citep{Newkirk:1961} was used to relate the observed emissions to the height of the emission source.

Figure \ref{fig1}, panels B and D, show type II burst dynamic spectra. Multiple bands are visible for the June 12 type II radio burst, indicating that this event is rather complex. Both fundamental and harmonic emissions were observed for that event, starting at 00:57:45 UT. The harmonic emission persisted longer, until about 01:07 UT, but was too faint to be measured. A strong type III burst was also observed at 00:53 UT---an indication of an impulsive release of electrons in the corona. We separated two emission lanes in each spectrogram and fitted the peak emission frequencies.

We performed the same analysis for the June 13 radio burst, which started at 05:38:13 UT. The fundamental emission was barely discernible in the radio spectrogram (panel D). However, there were two parallel bands of harmonic emission.


\section{RESULTS AND DISCUSSION}

Figure \ref{fig1}, panel E, shows particle flux enhancements possibly associated with the coronal shocks observed by AIA. The in-situ particle measurements were made by the Energetic and Relativistic Nuclei and Electron \citep[ERNE;][]{Torsti:1995} instrument on SOHO. The time series of energetic protons (between 1.68--90.5 MeV) show an impulsive flux increase at all energies on June 12, followed by an additional increase in the low energies on June 13. Vertical dashed lines denote onsets of the EUV waves. Below we investigate EUV wave observations and radio shock properties in an attempt to characterize the solar sources of the elevated particle fluxes.


\subsection{Kinematics of the EUV waves and radio shocks}
%1.1. KINEMATICS - VELOCITY, ACCELERATION - WHY IS THIS IMPORTANT?
Table \ref{table1} presents measured EUV wave front and radio shock kinematics. As described previously, we performed measurements of front edge positions along three linear profiles starting from the flare region (hereafter trials). These are denoted in roman numerals in the second column of the Table, together with the AIA channel. The third and fourth columns show initial velocities and acceleration, respectively, derived from second-order polynomial fits to the de-projected position measurements in two trials for each event. The rows in bold show trial averages for each wavelength, for each event.

For June 12, we obtained velocities of $\sim1275\pm44$~km~s$^{-1}$ for AIA/193 and $\sim1300\pm44$~km~s$^{-1}$ for AIA/211 channel. For June 13, we get $\sim731\pm22$~km~s$^{-1}$ for AIA/193 and $\sim741\pm31$~km~s$^{-1}$  for AIA/211 channel. The fits imply average decelerations of $-1000$~m~s$^{-2}$ for June 12 and $-800$~m~s$^{-2}$ for June 13.

\citet{Patsourakos:2010} studied the June 13 CME in EUV with AIA, between 05:34--05:43 UT. They fit circles to the expanding CME bubble and determined its kinematics. They found that the bubble front in the direction of propagation away from the solar limb accelerated to a maximum speed of ~400 km/s, after which it decelerated. They did not comment on the wave kinematics in that work.

\citet{Veronig:2010} studied a very similar dome-like CME and wave event off the eastern solar limb (seen from the STEREO-B spacecraft) with EUV observations They found upward expansion speeds of the dome-like wave of $\sim650$ km~s$^{-1}$. They also found that the EUV wavefront coincided with the white light transient observed by STEREO-B coronagraphs. This implies that the front edge of the white light emission may be caused by compressed elecron plasma behind the traveling shock. Quadrature position modeling was done by \citet{Patsourakos:2009b} to first show this connection.

Figure \ref{fig2} compares measured AIA EUV wavefront positions and estimated shock locations from radio observations. Wavefront positions measured in the 193 and 211 channels for the trial with lowest uncertainties are plotted as X-symbols. Diamonds denote radio shock positions. The June 12 radio emission occurred at lower heights than the EUV wavefront, suggesting electron acceleration away from the shock nose. Alternatively, the electron density model used might not apply for this case of open magnetic geometry (see Section~3.4). Radio emission started faster than the EUV wave, but decelerated completely by 01:00~UT, while the EUV wave continued to rise. By contrast, the radio emission on June 13 was split into two harmonic bands, which correlate very well with the EUV wave positions. This might imply local electron acceleration in front of and behind the nose of the traveling shock. Section~3.4 considers the coronal magnetic geometry in interpreting these observations.

%{\bf This might imply energetic electrons streaming along closed field lines low in the corona.} {\color{red} The previous statement is probably wrong. Usually it is assumed that the shock accelerates electrons and they go into its upstream region. Not very far from the shock front, an instability develops and we can see radio emission. That is why I'm not sure that it's worth speaking about closed field lines}

%The radio positions are systematically lower than the front, which might be due to our choice of coronal density model. If this effect is real, then the electron accelerations may occur lower in the corona, below the shock front. REFERENCES FOR ACCELERATION AT SHOCK FLANKS???




\subsection{Temperature and density behavior of the EUV waves}
%1.3. TEMPERATURE RESPONSE OF AIA CHANNELS - DEM from Mark;
%In the absence of those, take two limiting cases:
% - Assume only temperature changes(n=const). Discuss my attempts to reconstruct isothermal EM a la Schmelz et al. (2010)
% - Assume only density changes(T=const). Then the intensity is
% directly related to density, and the ratio of emissions at same
% wavelength but at different times gives the square of the density
% ratio.

To investigate the temporal and density properties of the EUV waves, we performed differential emission measure (DEM) analysis on the June 13 wave (we were not able to do so for the June 12 wave due to data constraints), using region-averaged pixel values in the six EUV Fe-dominated channels (94, 131, 171, 193, 211, 335~\AA). We hand-selected four regions (labeled R1-4 in top panel of Fig.\ref{fig3}) in two frames - 05:37:00~UT and 05:39:00~UT (hereafter T1 and T2) - corresponding to times before and during the wavefront passage. The first three regions were chosen to sample different parts of the wave; the fourth was chosen upstream of the wave for comparison. Calculations were done for 16 temperature bins between $logT=5.5-7.0$, following the Monte Carlo method as implemented in \citet{Weber:2004}. Results for regions 1 and 4 were not statistically significant, so our analysis was limited to times T1 and T2 in R2 and R3.
%using the SolarSoft code {\tt xrt\_dem\_iterative2.pro} with the AIA responses
% (see Table~\ref{table2})

For each region (of approximately 10,000 pixels), time, and channel, we constructed mean intensity data and errors.  The mean observation sets were then solved for their DEM distributions. The DEM solutions we quote provide model intensities with the smallest $\chi^2$ fit to the data, when folded through the AIA responses. (See bottom panels of Fig.~\ref{fig3}.) We considered the relative degree of model fits versus the difference between T1 and T2 data.)
%In Table~\ref{table2}, we present the $\chi^2$ values of the models to data and between the data. This illustrates that (a) the differences between the data for T1 and T2 are statistically significant, and (b) the models are very good fits to the data. (Because the errors are on the mean of approximately 10,000 pixels, they are extremely small compared to typical Poisson errors per pixel, and the absolute sizes of all $\chi^2$ values appear inflated.

The DEMs for regions 2 and 3 are plotted in the bottom panels of Fig.\ref{fig3}, where T1 is shown in red and T2 is shown in green. The Monte Carlo analysis produces multiple solutions by varying the data by the errors, and these are represented as clouds of colored dotted lines with a very small spread. It can be seen that observations for T1 and T2 are significantly different. We find that the DEM temperature profile does not change appreciably from T1 to T2, for either region, but the overall emission measure increases.
% The spread of solutions due to the errors is very small in these cases, such that they are mostly underneath the solid colored lines of the main solutions.

To roughly estimate the jump in density, we consider a simple model. Assume that all measured intensity is emitted along the region's line-of-sight only from the wave-affected volume, i.e., no foreground nor background emission. Also assume no change in temperature. Then, since the integrated DEM is the full emission measure (EM) of the volume, we may estimate the density ratio as:

\begin{equation}
\frac{n_{e2}}{n_{e1}} \sim \frac{\sqrt{EM_2}}{\sqrt{EM_1}} \sim 
  \frac{\sqrt{\int \mathrm{DEM}_2(T) \mathrm{d}T}}{
        \sqrt{\int \mathrm{DEM}_1(T) \mathrm{d}T}}
\end{equation}

For region 2, we find that $n_{e2}/n_{e1} \simeq 1.18$, and for region 3, we find that $n_{e2}/n_{e1} \simeq 1.12$, consistent with weak coronal shocks. For a more sophisticated model that accounts for foreground and background emission, the density changes within the wave-affected volume would have to be {\it even larger} in order to generate the observed change in intensities.  Therefore, we find that $n_{e2}/n_{e1} \simeq 1.12$ is a {\it lower} limit.




\subsection{EUV wave morphology}
%1.2. MORPHOLOGY OF THE EMISSION
%-Multiple ripples seen in AIA, radio type II - corresponding to acceleration in 'bays'?
%-Widening of the sheaths between the CME and the front.
% 
%Some questions:
%Why is the post-shock sheath not uniform in intensity (there seems to be brightening along the northern flanks)? What does this imply for the upstream plasma conditions? Higher intensity must mean more plasma heating/compression, so if the shock width and speed is the same, does that mean the Alfven speed is different, or perhaps that there is more pressure upstream?
%Why does the post-shock sheath change in width? Is there a connection between the sheath width, the sheath intensity, and the shock strength? Is this a function of position along the shock front, a function of time?

In both 211 and 193 channels, brightness increases downstream of the EUV wavefronts, relative to upstream. In the 211 channel this brightening is more pronounced, and is also more uniform throughout the downstream region. In the 193 channel, by contrast, the downstream material emits only close to the leading front. In both cases, ripples of emission behind the wavefronts expand as the fronts sweep through regions of upstream coronal plasma. These features persist until the transients leave the AIA FOV for both events. However, the downstream sheath emission on June 12 was dimmer than the emission on June 13 (where a CME bubble was seen).
%However, on June 12 the downstream emission brightness was reduced significantly in both wavelengths as the disturbance propagated radially outward to regions of lower density. This could suggest a lack of a driver behind the front, especially when contrasted with sheath emission behind the June 13 wavefront. There, a clear CME driver behind the front is seen.

\subsubsection{Lateral Overexpansion of the June 13 CME}

\citet{Patsourakos:2010} studied the CME of June 13, and found a strong lateral overexpansion of the CME bubble in the first ~4 minutes, after which the bubble expansion became equal in the radial and lateral directions (see top panel in their Fig.4.) Recently reported 3D numerical MHD simulations of coronal CME propagationin \citep{Das:2011} show that a pile-up compression (PUC) of coronal plasma may form between coronal shocks and the CMEs behind them. This occurred in the simulation whenever the CME expanded faster laterally than radially. Their interpretation is that as a CME expands fast laterally, plasma piles in front of it in a `sheath' behind the shock \citep{Opher:2010}. Additionaly, there was no significant temperature increase in the PUC in the simulations, consistent with our DEM results.

Comparing the intensities of the June 13 wave with results from \citet[their Fig.~4]{Patsourakos:2010}, we find the wave began increasing brightness significantly towards 05:38~UT, roughly coincident with the maximum speed of the CME. Even as the CME bubble aspect ratio reduced to $~$1, the overall wave brightness increased, peaking around 05:42~UT (after that the wave begins to disappear from the AIA FOV). Since the waves are quite dim, it was difficult to obtain quantitative observations. Future work will elucidate the temporal connection between lateral overexpansion and PUC formation. However, the DEM result of no significant temperature change support the modeling findings of a plasma compression sheath behind the shock from \citet{Das:2011}.


\subsection{Importance of the Magnetic Geometry for Particle Acceleration and Release}

Figure \ref{fig4} shows SDO/AIA (green) and STEREO-Ahead/EUVI \citep[red;][]{Howard:2008} difference images during both events, with a magnetic potential field source surface\citep[PFSS;][]{Schrijver:2003} model overlaid. On June 12 (left, top and bottom) the field geometry above the AR was very open, so particles were free to escape into interplanetary space as they gained energy. However, the complex magnetic topology does not allow us to address the possible sites of particle acceleration, and thus the discrepancy in positions and velocities between the EUV wave and radio shock.

%{\color{red}Large-scale loop systems extending south and east (bottom left panel) might explain the radio emission during that event - electrons accelerated in a shock expanding along the loop system could have caused that emission.}

On June 13 (Fig.~4, right panel), the magnetic geometry above the AR was much more closed - the shock might have been quasi-perpendicular at its nose, accelerating particles more effectively there (evidenced by the radio emission bands positioned in front of and behind the AIA wavefront). However, DEM analysis shows it was weak, so the 1 AU impulsive proton fluxes higher than 8~MeV did not increase above the already elevated levels (Fig.~\ref{fig1}, panel~E).


\section{SUMMARY}

We have presented observations of two western off-limb coronal waves in very high-cadence EUV imaging data. The waves were associated with metric type II bursts and significant increases in proton fluxes observed at 1 AU. We characterized the wave events in relation to the elevated particle fluxes at 1 AU. Our findings are:
1) The June 12 and 13, 2010 waves were large-scale, dome-like off-limb coronal transients, seen in EUV light. Enhanced emission sheaths followed the wavefronts.
2) The June 12 wave has a high initial speed ($\sim1287$~km~s$^{-1}$), but without a discernible driver, supported by a high average deceleration rate($\sim-1170$~m~s$^{-2}$).Similar behavior was observed in radio shock emission, although a discrepancy is clear between wave/shock positions and velocities. This might signify a more complex relation between the shock and wave, or alternatively, that the electron density model used for the radio data does not apply in this case. The June 13 wave started slower ($\sim736$~km~s$^{-1}$), but had a clear CME driver behind it sustaining its propagation, and consequently, a lower deceleration rate ($\sim-800$~m~s$^{-2}$).
3) DEM analysis of the June 13 wave event shows the enhanced emission was likely due to a density increase in the sheath behind the shock, and not to a temperature increase. We deduce from the emission measures ratio a density jump of at least $\sim1.12$.
4) EUV, radio, and in-situ observations, combined with a potential magnetic field model, reveal differences in the two events in terms of the possible field-to-shock orientation. In our interpretation, a more open field geometry of the June 12 event allowed protons accelerated impulsively (to $\sim50$~MeV) to escape quickly into interplanetary space. A closed field geometry during the June 13 event is supported by radio observations indicating the shock was effective in accelerating electrons at its nose, although proton fluxes above $\sim~8$~MeV at 1 AU did not increase appreciably.

The mechanisms of shock formation in the low corona are still under considerable debate. However, it seems that shocks do form low in the corona, and they are able to accelerate particles. The newly-introduced capability for multi-wavelength ultra-high cadence EUV observations of transients in the corona with SDO/AIA enables studying their dynamics in great detail. Based on our findings, the magnetic field geometry is important both for accelerating particles, and for their release into interplanetary space. Future work will involve analyzing multiple events and associated in-situ particle fluxes from multiple spacecraft, in order to constrain remote EUV wave observables significant for particle acceleration in the corona.% We will also perform modeling of particle transport between the Sun and 1 AU.

%However, in a study of type II bursts \cite{Gopalswamy:2005} determined that, on average, coronal shocks start between 1.6-2.2 $R_{s}$, suggesting that shocks do form low in the corona. We find very good associations of initial times, heights, and speeds between AIA observations and radio type II shock parameters. Follow-up work analyzing AIA and in-situ data will establish more definitively whether these transients are shocks. This work could impact particle acceleration theories and modeling, and improve significantly space weather predictions.


\acknowledgments
We acknowledge support under AIA subcontract SP02H1701R from Lockheed-Martin and NASA LWS EMMREM project NNX07AC14G. We thank David Long, Maher Dayeh, Marc De Rosa, Steve Saar, and Suli Ma for help and discussions.


%%\clearpage

%% Two options are available to the author for producing tables:  the
%% deluxetable environment provided by the AASTeX package or the LaTeX
%% table environment.  Use of deluxetable is preferred.
%%

%% Three table samples follow, two marked up in the deluxetable environment,
%% one marked up as a LaTeX table.

%% In this first example, note that the \tabletypesize{}
%% command has been used to reduce the font size of the table.
%% We also use the \rotate command to rotate the table to
%% landscape orientation since it is very wide even at the
%% reduced font size.
%%
%% Note also that the \label command needs to be placed
%% inside the \tablecaption.

%% This table also includes a table comment indicating that the full
%% version will be available in machine-readable format in the electronic
%% edition.

%\clearpage

%\begin{deluxetable}{cccccccc}
%\tabletypesize{\scriptsize}
%\rotate
%\tablecaption{Type II and III radio bursts temporally associated with AIA coronal limb waves\\June 12 and 13, 2010\label{table1}}
%\tablewidth{0pt}
%\tablehead{
%\colhead{Date} & \colhead{Station} & \colhead{Start time (UT)} & \colhead{End time (UT)} & \colhead{Type} & \colhead{High Freq. (MHz)} &
%\colhead{Low Freq. (MHz)} & \colhead{ESS\tablenotemark{a} (km/s)}}
%\startdata
%06/12 & CULG\tablenotemark{b} & 00:53 & 00:57 & III & 460 & 57X\tablenotemark{c} \\
%06/12 & HIRA & 00:54 & 00:56 & III & 1000 & 25X \\
%06/12 & PALE & 00:56 & 01:07 & II & 180 & 30 & \bf 3440 \\
%06/12 & LEAR & 00:57 & 01:07 & II & 179 & 29 & \bf 1360 \\
%06/12 & CULG & 00:58 & 01:09 & II & 170 & 57X & \bf 1300 \\
%06/13 & BLEN & 05:35 & 05:38 & III & 1550 & 1200X \\
%06/13 & CULG & 05:39 & 05:45 & II & 200 & 57X & \bf 700 \\
%06/13 & SVTO & 05:38 & 05:52 & II & 180 & 35 & \bf 621 \\
%06/13 & LEAR & 05:39 & 05:48 & II & 180 & 55 & \bf 665
%\enddata
%06/12 & CULG & 00:59 & 01:10 & II & 180 & 57X & \bf 600 \\
%% Text for table notes should follow after the \enddata but before
%% the \end{deluxetable}. Make sure there is at least one \tablenotemark
%% in the table for each \tablenotetext.
%\tablecomments{Radio burst information was taken from NGDC at \\             
%\url{ftp://ftp.ngdc.noaa.gov/STP/SOLAR\_DATA/SOLAR\_RADIO/SPECTRAL/2010/SPEC\_NEW.10} \\
%and Culgoora observatory website at 
%\url{http://www.ips.gov.au/Solar/2/6/1}}
%\tablenotetext{a}{Estimated Shock Speed}
%\tablenotetext{b}{The station codes are\\
%CULG - Culgoora observatory, Australia\\
%HIRA - Hiraiso observatory, Japan\\
%PALE - Palehua observatory, Hawaii, USA\\
%LEAR - Learmonth observatory, Australia\\
%BLEN - Blended observations from several stations\\
%SVTO - San Vito observatory, Italy}
%\tablenotetext{c}{X denotes that the emission reached the instrument's frequency limit.}
%\end{deluxetable}

\clearpage

\begin{deluxetable}{cccccc}

\tabletypesize{\scriptsize}
\tablecaption{EUV wave/radio shock kinematics, June 12 and 13, 2nd order fits\label{table1}}
\tablewidth{0pt}
\tablehead{
\colhead{Time} & \colhead{Channel/profile\tablenotemark{a} } & \colhead{Initial Velocity (km s$^{-1}$)} &
\colhead{Acceleration(km s$^{-2}$)} }

\startdata
06/12 00:56 & 193/II & 1169.34$\pm$29.31 & -0.88$\pm$0.16\\
06/12 00:56 & 193/III & 1381.59$\pm$32.70 & -1.29$\pm$0.17\\
06/12 00:56 & 211/II & 1180.25$\pm$31.67 & -0.95$\pm$0.17\\
06/12 00:56 & 211/III & 1418.08$\pm$31.56 & -1.55$\pm$0.17\\
\tableline
{\bf 06/12 00:56} & {\bf 193/AVG\tablenotemark{b}} & {\bf 1275.46$\pm$43.91} & {\bf -1.09$\pm$0.23}\\
{\bf 06/12 00:56} & {\bf 211/AVG} & {\bf 1299.16$\pm$44.71} & {\bf -1.25$\pm$0.24}\\
\tableline
\\
06/13 05:37 & 193/I & 774.29$\pm$18.90 & -0.97$\pm$0.18\\
06/13 05:37 & 193/II & 688.86$\pm$11.93 & -0.45$\pm$0.11\\
06/13 05:37 & 211/I & 791.89$\pm$25.30 & -1.15$\pm$0.24\\
06/13 05:37 & 211/II & 691.31$\pm$18.98 & -0.62$\pm$0.18\\
\tableline
{\bf 06/13 05:37} & {\bf 193/AVG} & {\bf 731.57$\pm$22.35} & {\bf -0.71$\pm$0.21}\\
{\bf 06/13 05:37} & {\bf 211/AVG} & {\bf 741.60$\pm$31.63} & {\bf -0.89$\pm$0.30}\\
\tableline
\\
\tableline
\tableline
06/12 00:57 & FUND & 2819.09$\pm$95.79 & -26.78$\pm$1.76\\
06/12 00:57 & HARM & 2905.60$\pm$147.92 & -46.79$\pm$5.30\\
06/13 05:39 & HARM & 589.54$\pm$27.21 & -0.53$\pm$0.23\\
06/13 05:38 & HARM & 610.93$\pm$14.97 & 0.28$\pm$0.11\\

\enddata
\tablecomments{Measurements for the 06/12 event started at 00:56 UT, for the 06/13 event - at 05:37 UT - the times we were able to first measure waves.}
\tablenotetext{a}{For radio measurements - emission type}
\tablenotetext{b}{Average of the profile measurements for that channel and event.}
\end{deluxetable}

%--------------------------------------------------------------------------------------------


%%  FIGURES
\clearpage

\begin{figure}
\epsscale{0.8}
\plotone{fig1.pdf}
\caption{Panels A and C: AIA/211 base difference images showing two stages of the June 12 and 13, 2010 coronal waves, respectively. The two frames in each panel are $\sim4$ minutes apart. Approximate positions of the wavefronts are in dashed black lines. Dotted lines in panel C outline the June 13 CME. The radial profiles along which velocity measurements were made are also shown. Panels B and D: radio spectra from Learmonth observatory for June 12 and 13, respectively. Panel B also shows a strong type III burst around 00:56~UT on June 12. Panel E: proton fluxes observed between June 11(DOY 162) and June 17(DOY 168), 2010 by the SOHO/ERNE instrument. Proton energies vary between 1.68-90.5 MeV. Vertical dashed lines show AIA waves onsets.\label{fig1}}
\end{figure}

\clearpage

\begin{figure}
\epsscale{0.8}
\plotone{fig2.pdf}
\caption{Time-height profiles of June 12 (top) and 13 (bottom) EUV waves and radio shocks.  Wavefront positions from the lowest uncertainty trial (see Table \ref{table1}) from AIA/193 and 211 channels are shown as X-symbols. Diamonds denote shock positions estimated from radio type II burst observations with the Newkirk density model.\label{fig2}}
\end{figure}
\clearpage


\begin{figure}
\epsscale{1.0}
\plotone{fig3.pdf}
\caption{Top - a snapshot of the June~13 event (base difference) with overlaid regions for which DEM solutions were attempted. Bottom - the DEM solutions for regions 2(left) and 3(right) as overlaid dotted histograms for the two times T1(red) and T2(green).\label{fig3}}
\end{figure}


\clearpage

\begin{figure}
\epsscale{1.0}
\plotone{fig4.pdf}
\caption{SDO/AIA (top panels) and STEREO-Ahead/EUVI (bottom panels) base difference images during the June 12 (left) and 13 (right) events. The PFSS model coronal fields are overlaid to show the topology in which the waves/shocks propagated.\label{fig4}}
\end{figure}

\clearpage
%\bibliography{bibliography}
\begin{thebibliography}{32}
\expandafter\ifx\csname natexlab\endcsname\relax\def\natexlab#1{#1}\fi

\bibitem[{{Bemporad} \& {Mancuso}(2010)}]{Bemporad:2010}
{Bemporad}, A., \& {Mancuso}, S. 2010, \apj, 720, 130

\bibitem[{{Das} {et~al.}(2011){Das}, {Opher}, {Evans}, {Loesch}, \&
  {Gombosi}}]{Das:2011}
{Das}, I., {Opher}, M., {Evans}, R., {Loesch}, C., \& {Gombosi}, T.~I. 2011,
  \apj, 729, 112

%\bibitem[{Domingo {et~al.}(1995)Domingo, Fleck, \& Poland}]{Domingo:1995}
%Domingo, V., Fleck, B., \& Poland, A.~I. 1995, Solar Physics, 162, 1,
%  10.1007/BF00733425

\bibitem[{Efron(1979)}]{Efron:1979}
Efron, B. 1979, The Annals of Statistics, 7, 1

%\bibitem[{{Evans} {et~al.}(2008){Evans}, {Opher}, {Manchester}, \&
%  {Gombosi}}]{Evans:2008}
%{Evans}, R.~M., {Opher}, M., {Manchester}, IV, W.~B., \& {Gombosi}, T.~I. 2008,
%  \apj, 687, 1355

\bibitem[{{Forbes} {et~al.}(2006){Forbes}, {Linker}, {Chen}, {Cid}, {K{\'o}ta},
  {Lee}, {Mann}, {Miki{\'c}}, {Potgieter}, {Schmidt}, {Siscoe}, {Vainio},
  {Antiochos}, \& {Riley}}]{Forbes:2006}
{Forbes}, T.~G., {et~al.} 2006, \ssr, 123, 251

\bibitem[{{Gallagher} \& {Long}(2010)}]{Gallagher:2010}
{Gallagher}, P.~T., \& {Long}, D.~M. 2010, \ssr, 135

\bibitem[{{Giacalone} \& {K{\'o}ta}(2006)}]{Giacalone:2006}
{Giacalone}, J., \& {K{\'o}ta}, J. 2006, \ssr, 124, 277

\bibitem[Howard et al.(2008)]{Howard:2008} Howard, R.~A., et al.\ 
2008, \ssr, 136, 67 

\bibitem[{{Maia} {et~al.}(2000){Maia}, {Pick}, {Vourlidas}, \&
  {Howard}}]{Maia:2000}
{Maia}, D., {Pick}, M., {Vourlidas}, A., \& {Howard}, R. 2000, \apjl, 528, L49

\bibitem[{{Mancuso} \& {Avetta}(2008)}]{Mancuso:2008}
{Mancuso}, S., \& {Avetta}, D. 2008, \apj, 677, 683

\bibitem[{{Mancuso} {et~al.}(2003){Mancuso}, {Raymond}, {Kohl}, {Ko}, {Uzzo},
  \& {Wu}}]{Mancuso:2003}
{Mancuso}, S., {Raymond}, J.~C., {Kohl}, J., {Ko}, Y., {Uzzo}, M., \& {Wu}, R.
  2003, \aap, 400, 347

\bibitem[{{Markwardt}(2009)}]{Markwardt:2009}
{Markwardt}, C.~B. 2009, in Astronomical Society of the Pacific Conference
  Series, Vol. 411, Astronomical Society of the Pacific Conference Series, ed.
  {D.~A.~Bohlender, D.~Durand, \& P.~Dowler}, 251--+

\bibitem[{{McLean} \& {Labrum}(1985)}]{McLean:1985}
{McLean}, D.~J., \& {Labrum}, N.~R. 1985, {Solar radiophysics: Studies of
  emission from the sun at metre wavelengths}

%\bibitem[{{M{\"u}ller-Mellin} {et~al.}(1995){M{\"u}ller-Mellin}, {Kunow},
%  {Flei{\ss}ner}, {Pehlke}, {Rode}, {R{\"o}schmann}, {Scharmberg}, {Sierks},
%  {Rusznyak}, {McKenna-Lawlor}, {Elendt}, {Sequeiros}, {Meziat}, {Sanchez},
%  {Medina}, {Del Peral}, {Witte}, {Marsden}, \& {Henrion}}]{Muller-Mellin:1995}
%{M{\"u}ller-Mellin}, R., {et~al.} 1995, \solphys, 162, 483

\bibitem[{{Newkirk}(1961)}]{Newkirk:1961}
{Newkirk}, Jr., G. 1961, \apj, 133, 983

\bibitem[{{Nindos} {et~al.}(2008){Nindos}, {Aurass}, {Klein}, \&
  {Trottet}}]{Nindos:2008}
{Nindos}, A., {Aurass}, H., {Klein}, K., \& {Trottet}, G. 2008, \solphys, 253,
  3

\bibitem[{{Opher}(2010)}]{Opher:2010}
{Opher}, M. 2010, {Shocks in heliophysics}, ed. {Schrijver, C.~J.~\& Siscoe,
  G.~L.} (Cambridge University Press), 193--+

\bibitem[{{Patsourakos} \& {Vourlidas}(2009)}]{Patsourakos:2009b}
{Patsourakos}, S., \& {Vourlidas}, A. 2009, \apjl, 700, L182

%\bibitem[Patsourakos et al.(2010a)]{Patsourakos:2010a} Patsourakos, S., Vourlidas, A., \& Kliem, B.\ 2010a, \aap, 522, A100 

\bibitem[{{Patsourakos} {et~al.}(2010){Patsourakos}, {Vourlidas}, \&
  {Stenborg}}]{Patsourakos:2010}
{Patsourakos}, S., {Vourlidas}, A., \& {Stenborg}, G. 2010, \apjl, 724, L188

\bibitem[{{Patsourakos} {et~al.}(2009){Patsourakos}, {Vourlidas}, {Wang},
  {Stenborg}, \& {Thernisien}}]{Patsourakos:2009a}
{Patsourakos}, S., {Vourlidas}, A., {Wang}, Y.~M., {Stenborg}, G., \&
  {Thernisien}, A. 2009, \solphys, 259, 49

\bibitem[{{Pick} {et~al.}(2006){Pick}, {Forbes}, {Mann}, {Cane}, {Chen},
  {Ciaravella}, {Cremades}, {Howard}, {Hudson}, {Klassen}, {Klein}, {Lee},
  {Linker}, {Maia}, {Mikic}, {Raymond}, {Reiner}, {Simnett}, {Srivastava},
  {Tripathi}, {Vainio}, {Vourlidas}, {Zhang}, {Zurbuchen}, {Sheeley}, \&
  {Marqu{\'e}}}]{Pick:2006}
{Pick}, M., {et~al.} 2006, \ssr, 123, 341

\bibitem[{{Raymond} {et~al.}(2000){Raymond}, {Thompson}, {St.~Cyr},
  {Gopalswamy}, {Kahler}, {Kaiser}, {Lara}, {Ciaravella}, {Romoli}, \&
  {O'Neal}}]{Raymond:2000}
{Raymond}, J.~C., {et~al.} 2000, \grl, 27, 1439

\bibitem[{{Reiner} {et~al.}(2007){Reiner}, {Kaiser}, \&
  {Bougeret}}]{Reiner:2007}
{Reiner}, M.~J., {Kaiser}, M.~L., \& {Bougeret}, J. 2007, \apj, 663, 1369

\bibitem[{{Reiner} {et~al.}(2003){Reiner}, {Vourlidas}, {Cyr}, {Burkepile},
  {Howard}, {Kaiser}, {Prestage}, \& {Bougeret}}]{Reiner:2003}
{Reiner}, M.~J., {Vourlidas}, A., {Cyr}, O.~C.~S., {Burkepile}, J.~T.,
  {Howard}, R.~A., {Kaiser}, M.~L., {Prestage}, N.~P., \& {Bougeret}, J. 2003,
  \apj, 590, 533

\bibitem[{{Roussev} {et~al.}(2004){Roussev}, {Sokolov}, {Forbes}, {Gombosi},
  {Lee}, \& {Sakai}}]{Roussev:2004}
{Roussev}, I.~I., {Sokolov}, I.~V., {Forbes}, T.~G., {Gombosi}, T.~I., {Lee},
  M.~A., \& {Sakai}, J.~I. 2004, \apjl, 605, L73

\bibitem[{{Schrijver} \& {De Rosa}(2003)}]{Schrijver:2003}
{Schrijver}, C.~J., \& {De Rosa}, M.~L. 2003, \solphys, 212, 165

\bibitem[{{Stone} \& {Tsurutani}(1985)}]{Stone:1985}
{Stone}, R.~G., \& {Tsurutani}, B.~T. 1985, Washington DC American Geophysical
  Union Geophysical Monograph Series, 34

\bibitem[{{Title} \& {AIA team}(2006)}]{Title:2006}
{Title}, A.~M., \& {AIA team}. 2006, in Bulletin of the American Astronomical
  Society, Vol.~38, Bulletin of the American Astronomical Society, 261--+

\bibitem[Torsti et al.(1995)]{Torsti:1995} Torsti, J., et al.\ 
1995, \solphys, 162, 505

\bibitem[Veronig et al.(2010)]{Veronig:2010} Veronig, A.~M., Muhr, 
N., Kienreich, I.~W., Temmer, M., \& Vr{\v s}nak, B.\ 2010, \apjl, 716, L57 

\bibitem[{{Vourlidas} {et~al.}(2003){Vourlidas}, {Wu}, {Wang}, {Subramanian},
  \& {Howard}}]{Vourlidas:2003}
{Vourlidas}, A., {Wu}, S.~T., {Wang}, A.~H., {Subramanian}, P., \& {Howard},
  R.~A. 2003, \apj, 598, 1392

\bibitem[{{Weber} {et~al.}(2004){Weber}, {De Luca}, {Golub}, \&
  {Sette}}]{Weber:2004}
{Weber}, M.~A., {De Luca}, E.~E., {Golub}, L., \& {Sette}, A.~L. 2004, 
  in IAU Symp.~223, Multi-Wavelength Investigations of Solar Activity, ed.\  
  {A.~V.~Stepanov, E.~E.~Benevolenskaya, \& A.~G.~Kosovichev} 
  (Cambridge: Cambridge Univ.\ Press), 321--328

\bibitem[{{Zank} {et~al.}(2008){Zank}, {Li}, \& {Verkhoglyadova}}]{Zank:2008}
{Zank}, G.~P., {Li}, G., \& {Verkhoglyadova}, O. 2008, in American Institute of
  Physics Conference Series, Vol. 1039, American Institute of Physics
  Conference Series, ed. {G.~Li, Q.~Hu, O.~Verkhoglyadova, G.~P.~Zank,
  R.~P.~Lin, \& J.~Luhmann }, 203--213


\end{thebibliography}


\end{document}

1406.2372.extracted_bib
[{{Bemporad \& {Mancuso(2010)]{Bemporad:2010 {Bemporad, A., \& {Mancuso, S. 2010, \apj, 720, 130
1406.2372.extracted_bib
[{{Das {et~al.(2011){Das, {Opher, {Evans, {Loesch, \& {Gombosi]{Das:2011 {Das, I., {Opher, M., {Evans, R., {Loesch, C., \& {Gombosi, T.~I. 2011, \apj, 729, 112 %
1406.2372.extracted_bib
[{Domingo {et~al.(1995)Domingo, Fleck, \& Poland]{Domingo:1995 %Domingo, V., Fleck, B., \& Poland, A.~I. 1995, Solar Physics, 162, 1, % 10.1007/BF00733425
1406.2372.extracted_bib
[{Efron(1979)]{Efron:1979 Efron, B. 1979, The Annals of Statistics, 7, 1 %
1406.2372.extracted_bib
[{{Evans {et~al.(2008){Evans, {Opher, {Manchester, \& % {Gombosi]{Evans:2008 %{Evans, R.~M., {Opher, M., {Manchester, IV, W.~B., \& {Gombosi, T.~I. 2008, % \apj, 687, 1355
1406.2372.extracted_bib
[{{Forbes {et~al.(2006){Forbes, {Linker, {Chen, {Cid, {K{\'ota, {Lee, {Mann, {Miki{\'c, {Potgieter, {Schmidt, {Siscoe, {Vainio, {Antiochos, \& {Riley]{Forbes:2006 {Forbes, T.~G., {et~al. 2006, \ssr, 123, 251
1406.2372.extracted_bib
[{{Gallagher \& {Long(2010)]{Gallagher:2010 {Gallagher, P.~T., \& {Long, D.~M. 2010, \ssr, 135
1406.2372.extracted_bib
[{{Giacalone \& {K{\'ota(2006)]{Giacalone:2006 {Giacalone, J., \& {K{\'ota, J. 2006, \ssr, 124, 277
1406.2372.extracted_bib
[Howard et al.(2008)]{Howard:2008 Howard, R.~A., et al.\ 2008, \ssr, 136, 67
1406.2372.extracted_bib
[{{Maia {et~al.(2000){Maia, {Pick, {Vourlidas, \& {Howard]{Maia:2000 {Maia, D., {Pick, M., {Vourlidas, A., \& {Howard, R. 2000, \apjl, 528, L49
1406.2372.extracted_bib
[{{Mancuso \& {Avetta(2008)]{Mancuso:2008 {Mancuso, S., \& {Avetta, D. 2008, \apj, 677, 683
1406.2372.extracted_bib
[{{Mancuso {et~al.(2003){Mancuso, {Raymond, {Kohl, {Ko, {Uzzo, \& {Wu]{Mancuso:2003 {Mancuso, S., {Raymond, J.~C., {Kohl, J., {Ko, Y., {Uzzo, M., \& {Wu, R. 2003, \aap, 400, 347
1406.2372.extracted_bib
[{{Markwardt(2009)]{Markwardt:2009 {Markwardt, C.~B. 2009, in Astronomical Society of the Pacific Conference Series, Vol. 411, Astronomical Society of the Pacific Conference Series, ed. {D.~A.~Bohlender, D.~Durand, \& P.~Dowler, 251--+
1406.2372.extracted_bib
[{{McLean \& {Labrum(1985)]{McLean:1985 {McLean, D.~J., \& {Labrum, N.~R. 1985, {Solar radiophysics: Studies of emission from the sun at metre wavelengths %
1406.2372.extracted_bib
[{{M{\"uller-Mellin {et~al.(1995){M{\"uller-Mellin, {Kunow, % {Flei{\ssner, {Pehlke, {Rode, {R{\"oschmann, {Scharmberg, {Sierks, % {Rusznyak, {McKenna-Lawlor, {Elendt, {Sequeiros, {Meziat, {Sanchez, % {Medina, {Del Peral, {Witte, {Marsden, \& {Henrion]{Muller-Mellin:1995 %{M{\"uller-Mellin, R., {et~al. 1995, \solphys, 162, 483
1406.2372.extracted_bib
[{{Newkirk(1961)]{Newkirk:1961 {Newkirk, Jr., G. 1961, \apj, 133, 983
1406.2372.extracted_bib
[{{Nindos {et~al.(2008){Nindos, {Aurass, {Klein, \& {Trottet]{Nindos:2008 {Nindos, A., {Aurass, H., {Klein, K., \& {Trottet, G. 2008, \solphys, 253, 3
1406.2372.extracted_bib
[{{Opher(2010)]{Opher:2010 {Opher, M. 2010, {Shocks in heliophysics, ed. {Schrijver, C.~J.~\& Siscoe, G.~L. (Cambridge University Press), 193--+
1406.2372.extracted_bib
[{{Patsourakos \& {Vourlidas(2009)]{Patsourakos:2009b {Patsourakos, S., \& {Vourlidas, A. 2009, \apjl, 700, L182 %
1406.2372.extracted_bib
[Patsourakos et al.(2010a)]{Patsourakos:2010a Patsourakos, S., Vourlidas, A., \& Kliem, B.\ 2010a, \aap, 522, A100
1406.2372.extracted_bib
[{{Patsourakos {et~al.(2010){Patsourakos, {Vourlidas, \& {Stenborg]{Patsourakos:2010 {Patsourakos, S., {Vourlidas, A., \& {Stenborg, G. 2010, \apjl, 724, L188
1406.2372.extracted_bib
[{{Patsourakos {et~al.(2009){Patsourakos, {Vourlidas, {Wang, {Stenborg, \& {Thernisien]{Patsourakos:2009a {Patsourakos, S., {Vourlidas, A., {Wang, Y.~M., {Stenborg, G., \& {Thernisien, A. 2009, \solphys, 259, 49
1406.2372.extracted_bib
[{{Pick {et~al.(2006){Pick, {Forbes, {Mann, {Cane, {Chen, {Ciaravella, {Cremades, {Howard, {Hudson, {Klassen, {Klein, {Lee, {Linker, {Maia, {Mikic, {Raymond, {Reiner, {Simnett, {Srivastava, {Tripathi, {Vainio, {Vourlidas, {Zhang, {Zurbuchen, {Sheeley, \& {Marqu{\'e]{Pick:2006 {Pick, M., {et~al. 2006, \ssr, 123, 341
1406.2372.extracted_bib
[{{Raymond {et~al.(2000){Raymond, {Thompson, {St.~Cyr, {Gopalswamy, {Kahler, {Kaiser, {Lara, {Ciaravella, {Romoli, \& {O'Neal]{Raymond:2000 {Raymond, J.~C., {et~al. 2000, \grl, 27, 1439
1406.2372.extracted_bib
[{{Reiner {et~al.(2007){Reiner, {Kaiser, \& {Bougeret]{Reiner:2007 {Reiner, M.~J., {Kaiser, M.~L., \& {Bougeret, J. 2007, \apj, 663, 1369
1406.2372.extracted_bib
[{{Reiner {et~al.(2003){Reiner, {Vourlidas, {Cyr, {Burkepile, {Howard, {Kaiser, {Prestage, \& {Bougeret]{Reiner:2003 {Reiner, M.~J., {Vourlidas, A., {Cyr, O.~C.~S., {Burkepile, J.~T., {Howard, R.~A., {Kaiser, M.~L., {Prestage, N.~P., \& {Bougeret, J. 2003, \apj, 590, 533
1406.2372.extracted_bib
[{{Roussev {et~al.(2004){Roussev, {Sokolov, {Forbes, {Gombosi, {Lee, \& {Sakai]{Roussev:2004 {Roussev, I.~I., {Sokolov, I.~V., {Forbes, T.~G., {Gombosi, T.~I., {Lee, M.~A., \& {Sakai, J.~I. 2004, \apjl, 605, L73
1406.2372.extracted_bib
[{{Schrijver \& {De Rosa(2003)]{Schrijver:2003 {Schrijver, C.~J., \& {De Rosa, M.~L. 2003, \solphys, 212, 165
1406.2372.extracted_bib
[{{Stone \& {Tsurutani(1985)]{Stone:1985 {Stone, R.~G., \& {Tsurutani, B.~T. 1985, Washington DC American Geophysical Union Geophysical Monograph Series, 34
1406.2372.extracted_bib
[{{Title \& {AIA team(2006)]{Title:2006 {Title, A.~M., \& {AIA team. 2006, in Bulletin of the American Astronomical Society, Vol.~38, Bulletin of the American Astronomical Society, 261--+
1406.2372.extracted_bib
[Torsti et al.(1995)]{Torsti:1995 Torsti, J., et al.\ 1995, \solphys, 162, 505
1406.2372.extracted_bib
[Veronig et al.(2010)]{Veronig:2010 Veronig, A.~M., Muhr, N., Kienreich, I.~W., Temmer, M., \& Vr{\v snak, B.\ 2010, \apjl, 716, L57
1406.2372.extracted_bib
[{{Vourlidas {et~al.(2003){Vourlidas, {Wu, {Wang, {Subramanian, \& {Howard]{Vourlidas:2003 {Vourlidas, A., {Wu, S.~T., {Wang, A.~H., {Subramanian, P., \& {Howard, R.~A. 2003, \apj, 598, 1392
1406.2372.extracted_bib
[{{Weber {et~al.(2004){Weber, {De Luca, {Golub, \& {Sette]{Weber:2004 {Weber, M.~A., {De Luca, E.~E., {Golub, L., \& {Sette, A.~L. 2004, in IAU Symp.~223, Multi-Wavelength Investigations of Solar Activity, ed.\ {A.~V.~Stepanov, E.~E.~Benevolenskaya, \& A.~G.~Kosovichev (Cambridge: Cambridge Univ.\ Press), 321--328
1406.2372.extracted_bib
[{{Zank {et~al.(2008){Zank, {Li, \& {Verkhoglyadova]{Zank:2008 {Zank, G.~P., {Li, G., \& {Verkhoglyadova, O. 2008, in American Institute of Physics Conference Series, Vol. 1039, American Institute of Physics Conference Series, ed. {G.~Li, Q.~Hu, O.~Verkhoglyadova, G.~P.~Zank, R.~P.~Lin, \& J.~Luhmann , 203--213