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Lecture 7: Radio Observations of Coronal Mass Ejections I

Hale COLLAborative Graduate Education (COLLAGE) Course 2017 Prof. Bin Chen (New Jersey Institute of Technology) Lectures 7-8 outline

• Radio preliminaries • Radiative transfer • Relevant emission mechanisms • Types of solar radio bursts This Lecture • Radio observations of CMEs • CME body • Thermal CME • Gyrosynchrotron CME • Type IV radio bursts • CME-driven shocks • White light/EUV imaging, UV spectroscopy, and signatures • type II radio bursts 1/16/2017 2 Radiation Fundamentals‣ Essential The surface containing can be any surface,Preliminaries real or imaginary; that is, it could be the physical surface of the detector, the source, or an imaginary surface anywhere along the ray. If energy from within the solid angle flows through the projected area in time and in a narrow frequency band of width , then • Specific intensity

(2.1) 1/16/2017 • Flux density2 Radiation Fundamentals‣ Essential Radio Astronomy Power is defined as the flow of energy per unit time, so the corresponding power is (2.9) • Units -2 -1 -1 • Flux �: ergs cm s Hz If the source angular size is , and• the1 Jy expression= 10-26 ergs cm for flux-2 densitys-1 Hz-1 is much simpler: • 1 solar flux unit (sfu) = 104 Jy • Specific intensity � : ergs cm-2 s-1 Hz-1 sr-1 Thus the quantitative definition of specific intensity or spectral brightness is • Sometimes radio images have units of Jy/beam (2.10) • Total flux: ergs cm-2 s-1 *CGS unit throughout this lecture This is usually the case for astronomical sources, and astronomers rarely use flux to describe sources so extended that the factor must be retained (e.g., the diffuse emission from our ). (2.2)

In practice, when should spectral brightness and when should flux density be used to describe a source? If a source is unresolved, meaning that it is much smaller in angular size than the point-source response of the eye or telescope observing it, its flux density can be and the MKSmeasured units butof its arespectral brightness cannot. To. the naked eye, the unresolved appears to be one of the brightest in the sky. Yet calling it a “bright star” is misleading because the total intensity of this relatively cool star is lower than the total Radio intensityastronomers of every almost hotter butalways more distantmeasure star frequencies,that is scarcely but visible observers to the eye. in Betelgeuse other wavebands appears “brighter” normally than measure most other wavelengths stars only rather than because it subtends a much larger solid angle and therefore its flux is higher. If a source is much larger than the point-source response, its frequencies,spectral so the brightness brightness at any per position unit wavelength on the source can be measured directly, but its flux density must be calculated by integrating the observed spectral brightnesses over the source solid angle. Consequently, flux densities are normally used to describe only relatively compact sources.

(2.3) is also widely used. The MKS units of are . The relation between the intensity per unit frequency and the intensity per unit wavelength can be derived from the requirement that the power in the frequency interval to must equal the power in the corresponding wavelength interval to :

(2.4) Figure 2.6: An illustration of the definition of flux density.

Thus The MKS units of flux density, , are much too big for practical astronomical use, so astronomers use smaller ones: http://www.cv.nrao.edu/~sransom/web/Ch2.html 7/49 http://www.cv.nrao.edu/~sransom/web/Ch2.html 3/49 Radiative Transfer

• In the absence of emission, absorption, or (the “free space”), the specific intensity � along a ray does not change. • However, if emission and absorption occurs, we use the radiative transfer equation �� = −� � + � �� -1 where � is the absorption coefficient (units cm ) -3 -1 and � is the emission coefficient (units ergs cm s Hz-1 sr-1). Radiative Transfer

• Defining the optical depth � = ��� (no unit) and the source function � = �/� the transfer equation can be written as: �� = � − � �� • For an isolated and homogeneous source the solution is � � = �(1 − � ) • When � ≫ 1, the source is optically thick and � ≈ � • When � ≪ 1, the source is optically thin and � ≈ �� Brightness Temperature

• While specific intensity can be expressed in units Planck Function of Jy/beam or SFU/beam, a simple and intuitive alternative is brightness temperature, which has units of Kelvin. Brightness Temperature

Note that at radio wavelengths

hn hn hv / kT << 1 ® ehn / kT -1»1+ -1= kT kT

The Planckian then simplifies to the Rayleigh-Jeans Law.

2hn 3 1 2n 2 B (T) = » kT n c2 ehn / kT -1 c2 It is useful to now introduce the concept of brightness

temperature TB, which is defined by 2n 2 I = B (T ) = kT n n B c2 B Rewriting the Radiative Transfer Equation

• Similarly, we define the as � = �� • Using our definitions of brightness temperature and effective temperature, the transfer equation can be rewritten

�� = −� + � ��

• Optically thick source, � ≫ 1, � ≈ � • Optically thin source, � ≪ 1, � ≈ �� Relevant Radio Emission Mechanisms: An Introduction

Bremsstrahlung Gyromagnetic Radiation Radiation

− − e− e e γ γ

Nucleus γ B Radiation

• Acceleration experienced in the Coulomb field − • At radio wavelengths, thermal e bremsstrahlung radiation is from γ virtually everywhere: quiet , active regions, flares, and CMEs Nucleus • Nonthermal bremsstrahlung is relevant to X-ray and gamma-ray emission from flares Gyromagnetic radiation

• Acceleration experienced in the • Gyroresonance radiation from thermal . Relevant in places with strong B field: e.g., active regions e− • Gyrosynchrotron radiation from relativistic γ electrons. Relevant when high energy electrons are present: e.g., flares and CMEs • gyrofrequency: one “natural frequency” of the solar corona B � = ≈ 2.8� MHz Plasma Radiation • , also known as “Langmuir wave”, occurs near the plasma frequency, another important natural frequency in the solar corona

�� � = ≈ 8980 � Hz �� • Plasma radiation arises when Langmuir waves are converted to (transverse) electromagnetic waves via wave-wave interactions. 1

29:129 – Plasma Oscillations— An application of electrostatics and classical mechanics

We consider a gas of electrons and positive (plasma). The plasma is overall neutral, i.e., the number density of the electrons and ions are the same. We think of the plasma as two interpenetratinglet’s do some derivation fluids — an negative electron fluid and a positive fluid.…

Plasma �

�

x x • Electric Coulomb force acts as the restoring force Schematic diagram of the plasma, with the Small volume of plasma in which the inset• ’s Law: showing a� typical� = 4�� volume. Integral form within ∯ the � ��� Electrons= 4�� are displaced to the right plasma with equal densities of positive ions by an amount x, while the ions are and electrons. fixed. • So � = = = 4��� �

Under normal conditions, there are always equal numbers positive ions and electrons in any • Newton’s 2nd law: � = −�� = −4�� ��, which has volume of the plasma, so the charge density = 0, and there is no large scale in the plasma. Now imagine that all of the electrons are displaced to the right by a small amount x, the form of a simple harmonic oscillator: while the positive ions are held fixed, as shown on the right side of the figure above. �̈ + �� = 0, The displacement of the electrons to the right leaves an excess of positive charge on the left side of the plasma slab and an excess of negative charge on the right side, as indicated by the dashedwhere rectangular� = boxes. The= positive2�� is the plasma frequency slab on the left and the negative slab on the right produce an electric field pointing toward the right that pulls the electrons back toward their original locations. However, the electric force on the electrons causes them to accelerate and gain , so they will overshoot their original positions. This situation is similar to a mass on a horizontal frictionless surface connected to a horizontal spring. The spring provides a restoring force that always acts to bring the mass back to its equilibrium position, thus producing simple harmonic motion. In the present problem, the electrons execute simple harmonic motion at a frequency that is called the electron plasma frequency, f pe pe 2 , where pe is the angular plasma frequency. We can apply what we know from electrostatics to compute the electric field that acts on the electrons, and use this in Newton’s second law to obtain pe . Let ne be the number density (in m3) of the electrons in the plasma, each having a charge of e. We can think of the positive region on the left side and the negative region on the right side as Plasma frequency: typical values

Listen to plasma oscillations when Voyager enters the local ISM

-3 Plasma � (cm ) � (MHz) Environment ISM 0.05 0.003 10 3 Low corona 10 900 Copper 10 2.8×10 Radio Observations in General

Thermal • Emission Mechanisms Non-Thermal

Chromosphere (~100 GHz) • Range of Observations Aurora (KHz) • Types of radio data Type III’s followed by Type II (45 mins)

+

Nobeyama Radioheliograph 17 GHz Solar radio bursts types from 1960s

• Type I: short duration, narrow band, non-drifting bursts. Origin unknown. • Type II: CME-driven shock (~1000 km/s) • Type III: fast electron beams (~0.3c), rapid frequency drift • Type IV: close magnetic structure -- CME body, post-CME reconfiguration • Type V: extended phase of type III Examples from the Green Bank Solar Radio Burst Spectrometer (GBSRBS) Green Bank Solar Radio Burst Spectrometer

• http://www.astro.umd.edu/~white/gb/ • Located in the Green Bank Radio Quiet Zone, operating in 18-70 MHz, 70-300 MHz, 300-1000 MHz Credit of following images: Stephen White Type III’s followed by Type II (45 mins) Type III burst: fast-drift electron beam (4 mins) Type II followed by Type IV (2 hours) Type V: extended phase of Type III (6 mins) The Astrophysical Journal,789:4(9pp),2014July1 Iwai et al.

AMATERAS LCP 20110116 The Astrophysical500 Journal,789:4(9pp),2014July1 Iwai et al. 10.0 AMATERAS LCP 20110116 8.0 400 500 10.0 8.0 400 6.0 300 6.0 300 4.0 4.0 dB from BG 200 2.0 dB from BG

Frequency (MHz) 2.0 200 0.0 01:00 01:30 02:00 02:30 03:00 03:30 04:00 04:30 Frequency (MHz) 0.0 Type I bursts (Time 3 hours(UT) ) 01:00 01:30 AMATERAS02:00 RCP02:30 2011012303:00 03:30 04:00 04:30 500 10.0 Time (UT) 8.0 400 AMATERAS RCP 201101236.0 300 4.0

500 dB from BG 10.0 200 2.0 Frequency (MHz) 0.0 8.0 400 01:00 01:30 02:00 02:30 03:00 03:30 04:00 04:30 Time (UT) AMATERAS RCP 20110126 6.0 300 ) 500 10.0 8.0 4.0 400 6.0 2.0 dB from BG 200 300 4.0 dB from BG Frequency (MHz) 200 2.0 0.0

Frequency (MHz 0.0 01:00 01:0001:3001:30 02:0002:00 02:3002:3003:00 03:3003:0004:00 04:3003:30 04:00 04:30 Time (UT)Time (UT) Figure 1. Radio dynamic spectra observed< 1s with AMATERAS on 2011 January 16 (LCP,RCP top), 23 (RCP, 20110126 middle), and 26 (RCP, bottom). BG: background.

) 500 AMATERAS Solar Radio Spectra 100 10.0 100 182 8.0 400 80 181 80

(SFU) 6.0 300180 60 60 4.0 179 dB from BG

Flux (SFU) Iwai et al. 2014

Frequency (MHz) 2.0

200178 Radio Intensity 40 Frequency (MHz 40 0.0 01:00177 01:30 02:00 02:30 03:00 03:30 04:00 04:30 20 0.0 0.1 0.2 0.3 0.4 0.5 0.6 0.7 Time20 (UT) Time (Sec) 0.0 0.1 0.2 0.3 0.4 0.5 0.6 0.7 Figure 1. Radio dynamic spectra observed with AMATERAS on 2011 January 16 (LCP,Time top),(Sec) 23 (RCP, middle), and 26 (RCP, bottom). BG: background. Figure 2. Left: an example of the radio dynamic spectra of a type-I burst element. The black diamond denotes the detected peaks of the type-I burst. The black asterisks are FWHMs of the duration and bandwidth of the burst. Right: the observed light curve of the type-I burst includes its peak flux (dotted line) and fitted Gaussian function of the observed light curve (solid line). The diamond indicates the burst peak and the triangles indicate the background continuum.SFU:solarflux unit (1 SFU 10 22 Wm 2 Hz 1). = − − − AMATERAS Solar Radio Spectra 100 The middle and bottom left panels of Figure 3 show the The100 correlation coefficient is 0.12. Hence, there is almost no distributions182 of the duration and bandwidth, respectively. The relationship. The middle and bottom right panels show a scatter solid and dotted lines show the power-law and exponential plot between the peak flux intensity and bandwidth and a scatter fitting curves, respectively. The distributions of the duration and plot between the duration and bandwidth, respectively. These bandwidth are more similar to exponential distributions than to figures also show small correlations. These results suggest that power-law181 distributions. the peak flux, duration,80 and bandwidth do not correlate with The top right panel of Figure 3 shows a scatter plot of the each80 other. In other words, there is no similarity in the shapes peak flux intensity versus the duration of the type-I bursts. of the burst (SFU) spectral structures. 180 3 60 60 179 Flux (SFU) Frequency (MHz)

178 Radio Intensity 40 40 177 20 0.0 0.1 0.2 0.3 0.4 0.5 0.6 0.7 20 Time (Sec) 0.0 0.1 0.2 0.3 0.4 0.5 0.6 0.7 Time (Sec) Figure 2. Left: an example of the radio dynamic spectra of a type-I burst element. The black diamond denotes the detected peaks of the type-I burst. The black asterisks are FWHMs of the duration and bandwidth of the burst. Right: the observed light curve of the type-I burst includes its peak flux (dotted line) and fitted Gaussian function of the observed light curve (solid line). The diamond indicates the burst peak and the triangles indicate the background continuum.SFU:solarflux unit (1 SFU 10 22 Wm 2 Hz 1). = − − −

The middle and bottom left panels of Figure 3 show the The correlation coefficient is 0.12. Hence, there is almost no distributions of the duration and bandwidth, respectively. The relationship. The middle and bottom right panels show a scatter solid and dotted lines show the power-law and exponential plot between the peak flux intensity and bandwidth and a scatter fitting curves, respectively. The distributions of the duration and plot between the duration and bandwidth, respectively. These bandwidth are more similar to exponential distributions than to figures also show small correlations. These results suggest that power-law distributions. the peak flux, duration, and bandwidth do not correlate with The top right panel of Figure 3 shows a scatter plot of the each other. In other words, there is no similarity in the shapes peak flux intensity versus the duration of the type-I bursts. of the burst spectral structures.

3 Other complex types

Strongest solar radio burst ever recorded, > 1 million sfu at ~1 GHz

From Dale Gary Decimetric radio bursts associated with flare termination shock

Eruption

Reconnection

Downflows

LT Radio (and HXR) Hot Loops

FP Radio Chen et al. 2015, Science 12

Radio bursts associated with propagating wave packets

Start Time 18:49:04 UT

1.6 1600 (a) 1530 MHz (b) 1307 MHz (d) 400

1500

350 1.4 1400 (c) 1198 MHz y (arcsec)

1300 Frequency (MHz)

300

1.2 1200

−950 −900 −850 1.5 3.5 5.5 x (arcsec) Duration (second) Wang, Chen & Gary 2017

Figure 3. (a)-(c) Three examples showing the morphology of the 12–25 keV HXR source (cyan), and the microwave source components (yellow, blue and red) obtained from an individual bursts in the dynamic spectrum (d). Yellow contours of (a)-(c) correspond to the emission ridge of the fiber burst, also indicated by yellow boxes in (d). Blue contours of (a)-(c) correspond to the adjacent absorption edge at the low frequency side of the corresponding emission ridge in (d). Red contours show the actual source hosting the fiber bursts, which are obtained by subtracting the emission (yellow) from the absorption (blue). The contour levels of radio sources are 30%, 70%, 95% and 99% of the their peak values. Around 18:50 UT when fiber bursts are very active, hot flare loops are only visible in the hot AIA passbands, as shown in the background images in (a)&(b), which correspond to AIA 131 A˚ and 94 A,˚ respectively. Around 19:28 UT, the coronal loops become evident in AIA 171 A,˚ as shown in the background of (c). itself, but also the background type IV continuum, which can be approximated by the intensity at the adjacent absorption edge. To obtain the source component purely due to the fiber bursts, we have applied a vector-based subtraction in the visibility space between the on-fiber emission and A nice illustration of 74 SOLAR AND RADIOPHYSICS

Quiet Sun

Gary & Hurford 2004

Figure 4.1. Characteristic radio frequencies for the solar . The upper-most charac- teristic frequency at a given frequency determines the dominant emission mechanism. The plot is meant to be schematic only, and is based on a model of temperature, density, and magnetic field as follows: The density is based on the VAL model B (Vernazza et al. 1981), extended to 105 km by requiring hydrostatic equilibrium, and then matched by a scale factor to agree with 5 the Saito et al. (1970) minimum corona model above that height (the factor of 5 was chosen to £ give 30 kHz as the 1 AU plasma frequency). Temperature was based on the VAL model to about 105 km, then extended to 2 106 K by a hydrostatic equilibrium model. The temperature is then £ taken to be constant to 1 AU. The magnetic field strength was taken to be the typical value for 1.5 active regions given by Dulk & McLean (1978), B =0.5(R/R 1) . For the ∫(øÆ =1) Ø °2 curve, a scale height L is needed. We used L = H0(T/T0)(R/R ) where H0 =0.1R and 6 Ø Ø T0 =2 10 K. Near the Sun, the curves apply to active regions. £

third harmonic ∫ =3∫B. Figure 1 shows that the 3∫B line lies above the øÆ =1 level down to 1–2 GHz, and extends up in frequency to 20 GHz—both of which agree well with the observed range. During bursts, gyroemissionª is more typically at ∫ = 10∫B, from which we see that gyroemission during bursts can extend to 800–900 MHz in the decimetric range. At mm wavelengths 100 GHz, bursts can be dominated by either free-free or gyroemission, dependingª on the number and energy of emitting particles. Outside of flares, the emission above 20 GHz is entirely due to free-free emission. Okay, let’s apply these to coronal mass ejections! Standard Flare-CME Model

Forbes-Lin model

From Lecture 5 (Prof. Jiong Qiu) CMEs in white light

From Angelos Vourlidas From Lecture 5 (Prof. Jiong Qiu) White light emission: Thomson scattering

• Not too far away from the solar disk, white light emission is dominated by Thomson scattering off free electrons (“K- corona”) • The corona is extremely optically thin, the white light brightness is � ≈ ��� = ����� v Thomson cross-section � = � ≈ 6.65×10cm v Source function � is related to the incident intensity from the and the viewing angle

See Prof. Steven Cranmer’s lecture 6 of COLLAGE 2016 for more details Review of CME White Light Observation

• White light emission is due Cavity to Thomson scattering, Front which goes as ��� • We can see a bright WL emission because: Core • It has more mass than the ambient Occulter • It is extended along the LOS Streamers • It is close to the plane of maximum scattering White Light CME

“Ice cream cone” model brightness

LASCO/C2 observed brightness

The “ice cream cone” model From Sarah Gibson Thermal radio emission from CME

• CME body also produces thermal bremsstrahlung radio emission • Coulomb “collisions” involve both ions and electrons − e • Radio intensity � ∝ ∫ ���� ∝ ��, known as “emission measure” γ • Radio intensity weakens much faster at larger distance, however relatively Nucleus more straightforward for modeling • Possible to see thermal radio CMEs against the disk (no occulter) Detectability of thermal radio CME

• Well, no problem for dense and cool prominences

� ∝ ���� ∝ �� Bremsstrahlung radiation / � ∝ �

NoRH 17 GHz 226 SOLAR AND SPACE WEATHER RADIOPHYSICS

CMEs have coronal temperatures and low densities and hence their free-free emission is expected to be very optically thin and difficult to observe, especially in the presence of the much brighter emissions from non-thermal mechanisms (Figure 11.2).

1.2.2 Nonthermal gyrosynchrotron. Gyrosynchrotron emission is routinely observed in flares (Bastian, Benz & Gary 1998) and could be present in CMEs, since they are capable of accelerating electrons to high energies (Kahler et al. 1986). The gyrosynchrotron emission from even a small num- ber of non-thermal electrons, entrained in the CME magnetic field, can easily exceed the thermal emission by a few orders of magnitude. Therefore, it is much easier to detect. Due to the dependence of gyrosynchrotron emission on the magnetic field, the brighter signal will correspond to the locations of the strongest magnetic field within the CME but the overall morphology will be similar to the white light CME. Gyroemission at low frequencies is suppressed in the presence of plasma where the index of deviates from unity Thermal radio detectability of CME body(Razin-Tsytovich suppression). This effect can be used as an additional diag-

• Contrast of CME/QS ~ 1:10 Simulated free-free radio spectra in meter & dm range, but much smaller at higher frequencies • Susceptible to be ”blinded out” by intense plasma radiation and/or nonthermal gyrosynchrotron radiation • Possible to detect with an instrument with large FoV and sufficient dynamic range

From Bastian & Gary 1997 Figure 11.2. Simulated free-free radio spectra for the quiet Sun and a typical CME (ne = 7 3 10 3.5 10 cm° , Te =2.5 MK, L =4.4 10 cm). The effects of refraction and reflection £ £ at the plasma layer are shown by the dashed and dotted lines, respectively. Diamonds are measurements by Zirin, Baumert & Hurford (1991). The bars, below 1 GHz, represent the range of measurements at a few relevant frequencies from Lantos et al. (1980) and Wang, Schmahl & Kundu (1987). From Bastian & Gary (1997). BASTIAN AND GARY: RADIO IMAGING OF CORONALSimulation from a toy CME model MASS EJECTIONS 14,035

1992, at 0233:46and 0259:22UT. Theseimages correspond to (a) Early CME (b) Late CME a long-durationevent studied by Tsunetaet al. [1992],an event which was accompaniedby a type II radio burst and may • This is from an ideal representthe SXR counterpartof a CME blowingout higher up in the corona [Tsuneta,1996; H. S. Hudson,private com- • In reality, detection using munication,1996]. By usingtwo separateYohkoh images, we this difference imaging includeintrinsic variations in solarfeatures that may havelittle technique would to do with the CME itself (e.g., solarrotation and variations inevitably suffer from within activeregions). Assuming a constanttemperature of 2.5 (c) Difference (d) Difference X 30 MK for the SXR-emittingplasma, the imageswere converted confusion due to to mapsof the columnemission measure. Assuming the radio uncleaned sidelobes emissionis due solely to thermal bremsstrahlungemission from the SXR-emitting plasma,the SXT maps of emission measureare easilyconverted to (frequency-dependent)maps of optical depthwhich are, in turn, convertedto mapsof the radio flux at selectedfrequencies. A backgrounddisk was added to each map; the contributionof the backgroundwas estimatedusing the simpletwo-component model of Zirin et al. Figure From Bastian & Gary 19973. The model brightnessdistribution used in the sim- [1991].Each map was convolvedwith a Gaussianwith a width ulationfor a referencefrequency of 960 MHz: (a) the "early" whichvaried as v: to simulatethe effects of coronalscattering. CME; (b) the "late" CME; (c) their difference(late minus The full-width at half maximum(FWHM) of the convolving early);(d) sameas Figure3c. The contourlevels are 100, 200, Gaussian at i GHz was 60". 500, 1000, 2000 x1000 K. Those of Figure 3c are _+1 times In reality, the solar corona is not isothermal.The coronal thoseused in Figures3a and 3b and thoseof Figures3d are 30 plasmais multithermaland is describedby a differentialemis- timessmaller than thoseof Figure 3c. The CME is modeledas sion measurewhich varies as a function of line of sight.The a hemisphericalshell of finite thicknessviewed from the side, and is not visible in the contours in a-c, but can be seen in effect of the presenceof multitemperatureplasma along the Figure 3d. line of sight is to reduce the radio brightnesstemperature somewhatat decimeterwavelengths [e.g., Klimchuk and Gary, 1996;Vourlidas and Bastian,1996, and reference therein] com- pared to that expectedfor a hot, isothermalplasma. While a flux. The normalized,expanded loop was again addedto the more realisticmodel of the plasmawould lead to quantitative model Sun, this time based on the SXT image obtained at differences,we believe the model brightnessdistributions are 0259:22 UT. We refer to this model as the "late" CME. The suf•cient for the determiningthe circumstancesunder which two model brightnessdistributions and their difference are CMEs may be detected. shownin Figure 3. Note that in the differenceimages there is 3.2.2. CME. For rough consistencywith the eruptive substantialresidual brightness present. This is dominatedby event identifiedin the SXT imageson February21, 1992,we the presenceof the SXR eruptive event toward the NE, but haveintroduced a modelCME at a solarelongation of • 2 Rs, there is a contributiondue to solar rotation during the course displacedradially from the eruptiveevent seen in soft X rays. of the 25 min between images. We note that CMEs may also originate in sourcesover the 3.3. Creation of Model Databases limb, in which caseearly changesin the radio brightnessdis- tributionwould be occultedby the limb. We discussthe detec- Model databaseswere constructedfor the strawmanarray tion of CMEs againstthe solar disk in section4.4. The CME configurationdescribed above. A modelbrightness distribution model is unrealisticto the extentthat it is modeledas a simple at a givenfrequency was Fourier transformed.The transform hemisphereof finite thickness(about 20% of the radius).Real wasthen sampled("degridded") on the u v pointscorrespond- CMEs would showfar more structure(e.g., bright legs and a ing to a frequencyof interestat a time correspondingto local tendencyto showthe three part structuredescribed in section noon at the array site. The procedurewas repeated for six 1). Nevertheless,we believethe model CME is approximately frequenciesbetween 600 and 1200 MHz. The resultswere the right scaleand contrastand is thereforesuf•cient to test combinedinto two multifrequencydatabases, one for the early various detection schemes.The electron number densityis CME and one for the late CME. From these databases we then assumedto be constantthroughout the shell. The densityis computedsnapshot maps at each discretefrequency and fre- determinedby the total massof the CME. For the exampleswe quencysynthesis maps from the combinationof frequencies. discusshere, we assumea totalmass of 10•6 gm,which gives a The instrumentresponse function (the dirty beam) was then numberdensity of a few x107cm -3 in the shell.A mapof the deconvolvedfrom the model brightnessdistributions as we column emissionmeasure of the shell was computed, con- now describe. verted to flux units and added to the model of the Sun based on the SXT image at 0233:46UT. This model is referred to as the 4. CME Detection Schemes "early" CME. To simulatesource expansion, we expandedthe CME in a 4.1. Direct Detection self-similarfashion by 20%. The volumeof the shellincreased The most straightforwardexercise is to see whether the by 1.23• 73%. The numberdensity decreased by a corre- model CME can be detectedthrough direct snapshotimaging spondingamount, but the columnlength increased by 20%, so at a singlefrequency, or by meansof frequencysynthesis [see, the columnemission measure changes by 1.2/(1.2)6 • -60%, e.g., BastJan,1989]. A complicationis that no joint spatial/ as does the optical depth, the brightness,and therefore the spectraldeconvolution algorithm is availableat present,al- A close-to-reality case

14,034 BASTIAN AND GARY: RADIO IMAGING 14,036OF CORONAL MASS EJECTIONSBASTIAN AND GARY: RADIO IMAGING OF CORONAL MASS EJECTIONS dt = 5 min dt = 10 min 37 antennas some bandwidthin order to improve samplingin the Fourier lOOO VLA as a means of removingthe preflare background.The domain [Conwayet al., 1990] and therefore reduce sidelobe scheme has worked well for bursts with a duration of tens of levels.In addition,we assumethat many short antennaspac- secondsobserved by the VLA, becausethe changein array ingsare neededfor goodsurface brightness sensitivity, while at geometryis small over this duration.However, the error made the same time the requirementfor an angular resolutionof by subtractingvisibility V(u •, v•) from [/'(/22,V2) increasesas 5OO order 40"/v9requires maximum baselines of 1.5 km, where v9is t 2 -- t• increases.We therefore investigatewhether such a the frequencyin gigahertz. subtractionwill work over the longer times over which CMEs We considera 37-elementarray (Figure 2). The arrayhas evolve,possibly tens of minutes. three linear arms arrangedin an inverted"Y" (120ø between arms), 12 elementsin eacharm, plus one in the middle.The To test the approximatedifferential detection scheme, we distancebetween antennas along each arm is initiallygiven by vector-subtractedvisibilities of the "early CME" from thoseof d =ab n, wheren = 0,...,ll, b = 1.5, anda = 3m. This the "late CME" for time differencesin array geometryof 5, 10, logarithmicspacing is chosenso that when frequencysynthesis 20, and 30 min. Note that the actual difference in time between techniquesare employed,sampling over a bandwidthratio of the two SXR imagesused in the modelsis about 25 min. The 1.5 allowscomplete spatial sampling in the radial directionout amountof solarvariability between two successivesnapshots is -5OO to the maximum baselinelength. The site latitude qbof the therefore exaggeratedfor time differencesless than 25 min. model instrumentwas chosento be qb= 35ø. In order to obtain Dirty maps and beams were made as describedabove and a circularinstrumental point spreadfunction when the Sun is deconvolvedusing identical parameters. The resultsare plot- at zero declination,the north-southspatial components of the ted as contoursof SNR in Figure 4. The residualsare much -1000 • , , , I , , i , I • , , , I , , , , antennalocations have been divided by cosqb; that is, the array better than those resultingfrom direct imaging,and now the - lOOO -500 o 500 ooo hasbeen stretched NS by a factor 1/cosqb • 1.22.The resulting CME can be seen. As noted above, the CLEAN box did not X (m) length of the east and west arms is roughly 818 m, while the include the CME. Sidelobes due to the CME were therefore Figure 2. The model array configurationused for both direct north arm is 942 m long. The maximumbaseline lengths are not CLEANed from the maps,and the rms residualis some- A hypothetical 37and approximatedifferential -element Ydetection schemes. -shaped arrayTwelve anten- about1.5 km, yieldingan instrumentalresolution of roughly84 nas are distributedlogarithmically along each arm. Figure 4.dt Frequency= 20 minsynthesis maps dt formed= 30 min using approxi- what higher as a result. The CME is seen for all values of arc secat 500 MHz (X = 60 cm), 28 arc secat 1.5 GHz (X = t 2 -- t• consideredbut with a SNR which decreaseswith 20mate cm), differentialand 2.8 arcdetection. sec at 15 GHzThe time(X = difference2 cm). Manybetween othersnap- shotsis (a) 5 min, (b) 10 min, (c) 20 min, and (d) 30 min. The increasingt 2 -- t• due to the changingarray geometry,as antennaSimulated difference imaging configurations are possible,of course.results at 960 MHz A studyto iden- contour levels are _+30, _+50, _+100-,_+15o, and _+200-.Note expected.As we approachtime differencesof 20-30 min, the instrumentmaybe used to detect and image CMEs with an tify an optimum configurationwill be needed once all science angular resolutionof-1 arc min. that the confusionlevel on the diskis increasinglyexaggerated CME is no longer readily detected at discretefrequencies objectivesfor time have differences been clearly less thanidentified. 25 min. The number and configurationof antennasemployed in above800 MHz (not shownin Figure4). In the averagemap it such an array is a complex optimization problem. A two- 3.2. Source Models can be seen,although the significanceof detectionis lessthan dimensionalarray of a large number of small antennasis de- We wish to detect thermal free-free emission from CMEs in 100-.Since we useda fairlylarge mass (10 •6 g) for ourmodel sirablein order to obtain the best possibleimage quality. On thethough decimetricsuch wavelength techniques range may in be the available presencein of thebright near and future CME, this is of some concern becausewe would like to detect the other hand, the minimum possiblenumber that achieves complex[Kommemission et al., 1997].on the In solar lieu disk. of a Thejoint emission deconvolutionfrom CMEsalgorithm CMEs lessmassive than the model CME crudelysimulated here. this goal is desiredin order to keep costslow. Other consid- is weextremely have mapped optically the thin, source as indicatedand deconvolved in Figurethe 1; instrumenttheir In Figure 5 we plot the rms residual in units of brightness erations include frequency coverage, frequency agility, and brightnessresponsetemperature function atis expecteddiscretefrequencies. to be of orderThe TB --deconvolved 102- temperaturefor the CLEAN maps shownin Figure 4. The temporalresolution. A generalpurpose solar-dedicated array, 10mapss K inare the thendecimetric averaged.range. Note The that brightness frequencytemperature synthesis onthrough dotted-dashedline representsthe rms residualbrightness A TB operatingin the frequencyrange of -300 MHz to 26.5 GHz is thea solar linear disk combinationvaries considerably of mapsdepending differs onfrom the afrequency joint spatial/ resultingfrom the direct detectionscheme, while the asterisks currentlyunder consideration[see Bastian and Gary, 1997].Its andspectral the typedeconvolution, of activity present.which Foris inherently the purposesnonlinear. of these representA TB resultingfrom approximatedifferential detec- purposeis to studya wide rangeof solarphenomena, including simulations,The mapswe wereassume deconvolvedthat flare-associatedusing the nonthermal Clark implementation emis- tion for the four time differencesconsidered (5, 10, 20, and 30 flares,filament eruptions, and CMEs, aswell as a hostof quiet sionsof havethe H6gbombeen largely CLEANexcluded algorithm.at the time We and used frequency 2000 iterations(or min). Owingto the fact that the residualbrightness on the disk Sun phenomena such as active regions, the quiet chromo- frequencies)with a loopin gain question. of 0.1We in allcan cases. then assumeThe samethat CLEAN the bright- box was is exaggeratedfor model time differencesof 5, 10, and 20 min sphereand corona,and coronalholes. It is within the context nessused temperaturein all cases. ofThe the rms disk residual rangeswas from computed chromospheric/for the clean (due to the fixed25 min differenceof the backgroundmodel), of this instrumentthat we wish to considerthe problem of transitionmap whichregion was values then used(104-105 to normalizeK) to coronalthe map.values The(few resulting CME detection,in part to determinehow the CME detection x106mapsK). therefore Hencea snapshotshow two-dimensional dynamic range ofdistributions a fewXl01-103 of bright- problem might drive the instrumentdesign. (seeFigure 1) will be neededto detectCMEs in the decimeter The current "strawman" design of the Solar Radio Tele- nessexpressed in terms of the signal-to-noise(SNR) ratio. wavelengthNote thatrange. the CLEAN box did not include the CME itself. We scope(SRT) callsfor •40 antennaswith diametersof 2-3 m It is important to use fairly realistic sourcemodels, at least find that the CME is not detectedby this scheme(SNR > 30-), and two or three large antennaswhich will be usedto calibrate insofaras confusingsources of emissionare concerned.Expe- neither in maps at discretefrequencies nor in the linear fre- the small antennasby referencingthem to cosmicflux stan- rience has shownthat "confusionnoise" can severelylimit the 2000 dards.The instrumentwill provide an imagingcapability be- dynamicquencyrange synthesis of Fourier(see Figuresynthesis 6).maps of the Sun [Bastian, tween 300 MHz and 26.5 GHz with a spectralresolution of a 1989].4.2. ConfusionApproximatenoise Differentialis due to the Detection presence of uncleanable few percent.Depending on the details of the observingpro- sidelobeclutter in the map, due to inadequateu v coverage gram, imagingspectroscopy could be performedover a broad and/orA timesecond variabilitydetection of thescheme source.is For approximate transientphenomena differential de- frequencyband on a timescaleas short as 1 s. onetection. mustrely The upon ideathe is to instantaneous vectorsubtract u v coveragethe complex("snapshot" visibilities at coverage)a time t•of from the array. thoseFor at time radio t 2 bursts, and thenthis tois not make a problem a map of the becauseresultthey and CLEAN are typically it usingfar thebrighter beam than at time the t background; 2. In otherwords, 3. Simulation of a Model Instrument and Data 0 10 20 confusionwe form noise a difference is thereforemap low. 12 However, * B 2 - forI• faint* B•, transient where I•, 2 Time difference 3.1. Array Configuration phenomena,representsconfusion the truenoise brightnesswill be adistribution major problem,at timesand itt•, is 2 and In choosingthe antennalocations, we assumethat frequency primarilyB •,2 representsthis issue which the we dirty address beamwith or the instrumental presentsimulations. point spread Figure 5. The rms residualof the linear frequencysynthesis synthesistechniques [BastJan, 1989] will be employedby the 3.2.1. Background brightness distribution. The back- mapsplotted as a functionof time differencebetween snapshot functionat timest•, 2. The asterisksymbol denotes a convolu- imagesfor variousdetection schemes. The dotted-dashedline SRT. Frequencysynthesis refers to a variant of aperture syn- ground radio brightnessdistribution was simulatedby using tion. Clearly,a simpleconvolution between the differencemap representsthe residuallevel of the direct detectionscheme thesiswhere several discrete frequenciesare sampledover two full-disk Yohkoh SXT imagesobtained on February 21, and a single instrumentalpoint spread function no longer describedin the text. The diamondsymbol represents the rms existssince B • 4: B 2. An approximatedeconvolution can be residual resultingfrom a 19-element array with a 19-element performed,however, using B 2. This techniquehas been em- redundantarray; the squaresymbol represents the samefor a ployed in the analysisof microwavebursts observed by the 37-element array with a 37-elementredundant array. 1992ApJ...390L..37G Some (rare) examples of thermal radio CMEs

• Gopalswamy & Kundu 1992. Observation made in 1986 using Clark Lake Radioheliograph at 73.8 MHz. v �_ ≈ 1×10 K v � ≈ 9.5 × 10 cm v � ≈ 2.7 × 10 g L164 RAMESH, KATHIRAVAN, & SASTRY Vol. 591

Fig. 1.—Difference image (05:54–05:06 UT) of the CME event observed Fig. 3.—Difference image obtained by subtracting the radioheliogram ob- with the LASCO C2 coronagraph on 2000 November 24. The inner circle 6 tained at 04:55 UT from 05:05 UT. The peakTb is 6.08 # 10 K, and the indicates the solar limb, and the outer circle is the occulting disk of the 6 contour interval is0.34 # 10 K. ≈ coronagraph. It extends approximately up to 2.2 R, from the center of the Sun. Solar north is straight up, and east is to the left. The CME can be clearly noticed as a bright structure above the northwest quadrant of the occulting sources on the disk, one can clearly observe enhanced radio disk. emission in close spatial correspondence with the white-light CME described above. Its estimated peak brightness temperature 5 northwest quadrant of the occulting disk of the coronagraph. ()wasfoundtobeTb 10 K. A comparison of the LASCO and ∼ Figure 2 shows the radioheliogram obtained with the GRH at GRH difference images clearly indicates that the radio enhance- 04:55 UT, the same day. Figures 3, 4, and 5 show the radio ment moved in the same direction as the white-light CME. Also, difference images corresponding to 05:05, 05:15, and 05:25 UT, their appearances are closely similar. We therefore conclude that with respect to the 04:55 UT image. In addition to the discrete the former is the radio counterpart of the latter. We estimated the velocity of the “radio CME” from the displacement of its Another examplecentroid in the difference images (Figs. 3, 4, and 5), and the

•L164White light CME RAMESH, KATHIRAVAN, & SASTRY Vol. 591

No. 2, 2003 EVOLUTION OF “HALO” CME CLOSE TO SUN L165

the CME event described here was not accompanied by any nonthermal continuum emission in the metric range, unlike the event reported by Bastian et al. (2001; see Sol.-Geophys. Data, 1998 June, for details). Therefore, it is possible that the CME- associated enhanced radio emission observed by us in the pre- sent case is most likely thermal in nature, and the excess emis- sion observed off the limb in the northwest quadrant of the Fig. 2.—Radioheliogram obtained with the GRH on 2000 November 24 GRH images is due to bremsstrahlung from the extra electrons 6 around 04:55 UT. The peakTb is 4.18 # 10 K, and the contour interval is associated with the halo CME. Its mass (M) is given by 0.27 # 106 K. The open circle at the≈ center is the solar limb. The instrument Fig. 4.—Same as in Fig. 3, but the timings are 04:55 and 05:15 UT. The 66 beam is shown near the bottom right corner. peakTb is 6.1 # 10K, and the contour interval is 0.39 # 10 K. Ϫ24 2 1/2 Ϫ1 1/2 ≈ M p 2 # 10 (5fTeb TL ) V g, (1)

wheref (MHz) is the observing frequency andL (R,) is the depth of the radio enhancement along the line of sight. The latter is unknown and is taken to be the same as the observed radial Fig. 1.—Difference image (05:54–05:06 UT) of the CME event observed Fig. 3.—Difference image obtained by subtracting the radioheliogram ob- width. The volume (V)oftheregionofenhancedemissionwas with the LASCO C2 coronagraph on 2000 November 24. The inner circle 6 tained at 04:55 UT from 05:05 UT. The peakTb is 6.08 # 10 K, and the indicates the solar limb, and the outer circle is the occulting disk of the 6 determined by multiplying its radial and lateral width with the Ramesh, Kathiravan & Sastry 2003, contour interval is0.34 # 10 K. ≈ coronagraph. It extends approximately up to 2.2 R, from the center of the depth along the line of sight. We assumed that the coronal plasma from Sun. SolarGauribidanur north is straight up, and eastRadioheliograph is to the left. The CME can be clearly in is a fully ionized gas of normal solar composition (90% noticed as a bright structure above the northwest quadrant of the occulting sources on the disk, one can clearly observe enhanced radio and 10% by number), and each electron is associated disk. emission in close spatial correspondence with the white-light with approximately2 # 10Ϫ24 g of material. In addition to the India at 109 MHz CME described above. Its estimated peak brightness temperature 5 above, we also derived the magnetic field strength (B)associated northwest quadrant of the occulting disk of the coronagraph. ()wasfoundtobeTb 10 K. A comparison of the LASCO and ∼ with the density enhancement assuming that the plasma b Figure 2 shows the radioheliogram obtained with the GRH at GRH difference images clearly indicates that the radio enhance- 0.05,asfoundbyVourlidasetal.(2000)forsomeoftheLASCO≈ 04:55 UT, the same day. Figures 3, 4, and 5 show the radio ment moved in the same directionFig. as the5.—Same white-light as in Figs. CME. 3 and Also, 4, but the timings are 04:55 and 05:25 UT. CMEs at about the same height range as the radio CME described The peakT is 5.6 # 1066K, and the contour interval is 0.29 # 10 K. difference images corresponding to 05:05, 05:15, and 05:25 UT, their appearances are closely similar. Web therefore≈ conclude that here. Table 1 lists the values of the different parameters of the with respect to the 04:55 UT image. In addition to the discrete the former is the radio counterpart of the latter. We estimated Ϫ1 latter mentioned above. It is well known that the distance (s) ,km s during the interval traveled by a CME in a given time interval (t)canbefound 73 עand 1185 of its 73 עthe velocity of the “radio CME”values from are the1087 displacement centroid in the difference images05:05–05:15 (Figs. 3, UT4, and and 5), 05:15–05:25 and the UT, respectively. This gives Ϫ2 since the initial speed (u)andtheacceleration(a)areknown. ms . Ϫ1 Ϫ2 121 ע it an effective acceleration of 157 ∼ For the values ofu p 1087 km s ,msa p 157 ,andt p 29 minutes (time difference between the last and first height measurement using GRH [05:25 UT] and LASCO data [05:54 3. ANALYSIS AND RESULTS UT], respectively), we found that the centroid of the radio CME should be located at a height of 5.78 R from the center of the Single frequency observations of CME-associated discrete ra- , Sun, at 05:54 UT. According to the LASCO measurements, the dio sources traveling outward to large heights ( 2–3 R )inthe , leading edge of the CME was located at a height of 5.71 R at solar corona are generally attributed to nonthermal∼ continuum , 05:54 UT. emission from moving type IV radio bursts (see Dulk 1980 and 7 references therein), since the observedTb ( 10 K; Wagner et al. 1981; Stewart et al. 1982; Gopalswamy≥ & Kundu 1989) is 4. CONCLUSIONS usually higher than that because of the emission from the back- We studied the kinematics of the halo CME event that took ground “quiet” Sun ( 106 K), which is purely thermal in nature. place on 2000 November 24 around 05:30 UT using the data ∼ But in the present case, theTb of the enhanced radio emission obtained with the GRH, close to the Sun. The speed of the observed at the location of the white-light CME is less than the CME in the low corona was estimated independently from the 6 electron temperature (Te ) of the solar corona ( 1.4 # 10 K; observed displacement of the associated thermal radio en- km sϪ1. This 73 ע Fludra et al. 1999). Also, no type IV emission≈ was reported hancement, and the average value is 1136 during our observing period on 2000 November 24 (Sol.- agrees well with the corresponding≈ speed estimated using Geophys. Data, 2001 January). We would like to point out here white-light observations. There is a good agreement between that optically thin from the nonthermal the extrapolated location of the radio CME and the height-time electrons entrained in the magnetic field of the CME could also measurements of the leading edge of the white-light CME ob- 4 5 give rise to low values ofTb ( 10 –10 K), as shown recently tained using the LASCO data, suggesting that the former cor- by Bastian et al. (2001) for the∼ event of 1998 April 20. But responds to the frontal loop of the white-light CME. The ac- Ϫ2 ms )estimatedby 121 ע Fig. 2.—Radioheliogram obtained with the GRH on 2000 November 24 we could not verify the above (through spectral index esti- celeration of the radio CME ( 157 6 ≈ around 04:55 UT. The peakTb is 4.18 # 10 K, and the contour interval is mation) in the present case, since the radio imaging data is us is consistent with the result published by Zhang et al. (2001), 0.27 # 106 K. The open circle at the≈ center is the solar limb. The instrument Fig. 4.—Same as in Fig. 3, but theavailable timings are at 04:55 only and one 05:15 frequency. UT. The However, it is to be noted that according to whom the CMEs undergo an acceleration of beam is shown near the bottom right corner. peakT is 6.1 # 1066K, and the contour interval is 0.39 # 10 K. b ≈ TABLE 1 Characteristics of the “Radio CME”

Time Coordinates of Centroid Position Anglea Brightness Temperature Electron Density Mass Magnetic Field Ϫ3 (UT) (R,) (deg) (K) (cm ) (g) (G) 05:05 ...... 0.22N, 1.06W 282 3.34 # 105 1.24 # 108 2.68 # 1016 2.40 05:15 ...... 1.09N, 1.41W 308 2.86 # 105 4.46 # 107 3.82 # 1016 1.33 05:25 ...... 2.03N, 1.81W 318 1.92 # 105 2.81 # 107 4.42 # 1016 0.86 a Measured counterclockwise from the solar north.