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Advances in Space Research 35 (2005) 739–754

Solar radio emissions

J.-P. Raulin *, A.A. Pacini 1

Centro de Radioastronomia e Astrofisica Mackenzie, Universidade Presbiteriana Mackenzie, Rua da Consolacao 896, Sao Paulo, 01302-907 SP, Brazil

Received 3 June 2004; received in revised form 1 March 2005; accepted 1 March 2005

Abstract

In this paper, we present a tutorial review which was presented at the first Advanced School on Space Environment (ASSE 2004).

We first describe the basics of radioastronomy definitions, and discuss radiation processes relevant to solar radio emissions like

plasma emission, free–free bremsstrahlung and gyromagnetic emissions. We illustrate these fundamentals by describing recent solar

radio observations and the constraints they bring on different solar physical parameters. We focus on solar radio emissions from the

quiet sun, active regions and during explosive events known as solar flares, and how the latter can bring quantitative informations

on the particles responsible for the emission. Finally, particular attention is paid to new radio diagnostics obtained at very high fre-

quencies in the millimeter/submillimeter range, as well as to radio emissions relevant to Space Weather studies.

� 2005 COSPAR. Published by Elsevier Ltd. All rights reserved.

Keywords: Solar radio astronomy; Radiation mechanisms; Solar flares; Submillimeter waves

1. Introduction

1.1. Why do we study the Sun at radio wavelengths?

Although the Sun is an average (size, temperature

and brightness) middle-aged star, its proximity to Earthof about 150,000,000 km (1 AU) makes it an extraordi-

nary laboratory for astrophysics. This proximity results

in a very strong signal in almost all the electromagnetic

spectrum, which allows to study the Sun with high sen-

sitivity, time and spectral resolution. The Sun�s large

angular size of �32 0 as seen from the Earth also permits

to describe in great details the solar surface and atmo-

spheric features. As a consequence, a wealth of phenom-ena has been found to occur through the Sun�satmosphere. For example, the well-known solar flares

are sudden releases of a great amount of energy within

high magnetic field regions called solar active regions.

0273-1177/$30 � 2005 COSPAR. Published by Elsevier Ltd. All rights reser

doi:10.1016/j.asr.2005.03.138

* Corresponding author.

E-mail address: [email protected] (J.-P. Raulin).1 Present address: Instituto Nacional de Pesquisas Espaciais, Sao

Jose dos Campos, Brazil.

During flares, energy stored in magnetic fields within ac-

tive centers is rapidly converted into thermal, kinetic

and mechanical energies. As a consequence, the local

plasma is heated to several or several tens of millions de-

grees, while other ambient particles are accelerated up to

high energies. During large flares as much as few 1038

electrons per second are accelerated, as deduced from

non-thermal bremsstrahlung hard X-rays, which pro-

vide a direct diagnostic of the electrons which collide

in the dense low solar atmosphere. Since a flare may last

few minutes, and assuming a typical coronal magnetic

structure filled with 1037 thermal particles, we see that

the accelerated electron production rate is very high. It

is also difficult to know on the relative positions of theacceleration regions (where the particles are accelerated)

and the radiation sites (where they emit). This raises the

question about the effects of the transport of these par-

ticles, on the observed time flare history, on the energy

spectrum of the particles, if, for example acceleration

and radiation sites are different. Solar flare radio obser-

vations give valuable informations in order to answer

these questions. As it gives also access to the highestenergy particles accelerated during flares, since high

ved.

mailto:[email protected]

740 J.-P. Raulin, A.A. Pacini / Advances in Space Research 35 (2005) 739–754

frequency radio diagnostics are more sensitive to these

particles than are hard X-rays produced by non-thermal

bremsstrahlung. Clues on the above mentioned ques-

tions: particle number, energy and spectrum, transport

of particles, are important to ultimately confront with

the existing acceleration models (see Miller et al., 1997,for a review on models and how they are constrained

by observations).

During solar flares, large scale magnetic field struc-

tures can be destabilized, and be propelled into the

interplanetary medium, along with the large masses it

contains, to form the so-called coronal mass ejections

(CMEs). It is now recognized that CMEs are the prin-

cipal drivers of the Space Weather and the near-Earthconditions. All these phenomena will emit at radio

wavelengths in a frequency range covering over seven

orders of magnitude from few tens of kHz up to few

tens or hundreds of GHz. Although many of the above

phenomena are well documented by the available

observations, they raise questions that are still unre-

solved. For example, it is not clear at all what trigger

the initiation of large CMEs. Plasma flows producedby filaments disruptions have been proposed (Wu

et al., 2000), although they might not be a general

cause for all CMEs (Simnett, 2000). The relationship

between flares and CMEs is still a subject of discussion

since many CMEs are detected without any flare man-

ifestation. Although CMEs occurrence follows approx-

imately the solar activity cycle, no clear associations

with the sunspot number have been found so far (How-ard et al., 1985; Webb, 2000).

In the remaining part of this section, we briefly de-

scribe how solar radio waves are observed. In the fol-

lowing section, are presented basic definitions and

relations useful in radio astronomy. Sections 3 and 4

are relative to radio emissions due to coherent and inco-

herent processes, respectively. Some aspects of the radi-

Fig. 1. Opacity of the Earth atmosphere as a fu

ative transfer of radio waves are discussed in Section 5,

and in Section 6 we present selected solar radio observa-

tions to illustrate the content of the previous sections.

Section 7 reviews and discuss some very high radio fre-

quency solar observations, and in Section 8 we present

radio observations relevant to Space Weather studies.In Section 9 we present our conclusions.

1.2. How do we observe solar radio waves?

Radio astronomy developed during the last fifty

years, that is relatively early compared to X-ray, ultravi-

olet and infrared techniques. The reason lies in the fact

that most radio radiation penetrates down to the groundlevel. In Fig. 1, we show how the Earth atmosphere af-

fects the incoming radiation from space as a function of

its wavelength. We see that the atmosphere is completely

transparent for waves between few centimeters and

10 m. Below 0.5 mm the radiation is almost totally ab-

sorbed. In the range 0.05–3 cm, there are some absorp-

tion regions mainly due to the presence of oxygen and

water vapour (Ulich, 1980), separated by windows ofvariable transparency. At the other end of the radio

spectrum, radiation with wavelengths greater than

about 20 m are totally blocked by the Earth�sionosphere.

To collect radio photons and concentrate them,

radiotelescopes are used. These instruments can be

single-dish antenna, as for example the 100 m diame-

ter Effelsberg radiotelescope shown in Fig. 2 and lo-cated in Germany. The main limitation of such

single-dish instruments are their poor capability of

resolving two distinct point sources on the sky, which

is given by the angular separation a (in radians)

a ¼ 1.22 � kD; ð1Þ

nction of wavelength (from NASA/IPAC).

Fig. 2. The 100 m diameter single-dish antenna radio telescope at

Effelsberg (Germany). After aries.phys.yorku.ca/bartel/SNmovie/

Effelsberg.jp.

J.-P. Raulin, A.A. Pacini / Advances in Space Research 35 (2005) 739–754 741

where k is the radio wavelength of observation and D is

the diameter of the antenna. Thus, even observing the

Sun at a wavelength of 3 cm, the radiotelescope shown

in Fig. 2 will not be able to resolve solar features with

angular separation smaller than 1 0, i.e., a projected size

on the solar surface of �42,000 km. Eq. (1) informs us

that higher spatial resolution can be obtained by usinglarger dishes, but clearly implying construction difficul-

ties as well as high costs. To overcome this problem

techniques have been developed based on the interfer-

ometry theory. The idea is to combine the signals from

many radiotelescopes arranged as to form a linear array

for example. The angular resolution of such an array is

given by an equation similar as Eq. (1), where D is not

anymore the diameter of a single antenna, but thegreater distance between two individual antennae. The

obtained array (or arrays) of antennae is called a radio

interferometer as shown for example in Fig. 3. This ar-

ray named the Very Large Array (VLA) is located near

Socorro (New Mexico, USA), and is composed of 27

antennae of 25 m diameter. Although the VLA is shown

in Fig. 3 in a compact configuration, the highest distance

between two individual antennas can reach �40 km.Thus, at a frequency of 5 GHz (k = 6 cm), the VLA

can resolve solar features with angular size less than

100, i.e., a projected size on the solar surface lower than

700 km. This kind of instrument will, thus, be able to

produce radio images of the Sun with enough angular

resolution to show details of the intensity distribution

on the solar disk. For further details on interferometric

Fig. 3. The very large array (Socorro, NM, USA) in its more compact

configuration Photo courtesy AUI/NRAO after http://sky-

server.sdss.org/dr1/en/proj/advanced/quasars/images/vla.gif.

techniques, the reader is referred to Kraus (1966), Chris-

tiansen and Hogbom (1985), McLean and Labrum

(1985).

Since the solar spectrum varies in time, and since the

spectrum of the radio emission can inform us about the

emission mechanism, it is common to record the solaremission as a function of both frequency and time. Such

way of representing the intensity of the radio emission is

called a dynamic spectrum, shown for example in Fig. 4.

2. Radio astronomy: basic relations

The radio emission received by a radio telescope ismeasured in terms of its intensity Im, which is the power

received per square meter, per frequency unit Hz, and

per solid angle unit (Kraus, 1966). Integrating Im over

the solid angle subtended by the emitting source we

get a quantity, Sm called the spectral flux density

Sm ¼Z

Im dX. ð2Þ

The unit of 1 W m�2 Hz�1 is quite large for most of theradio emitting astrophysical objects, and a new flux den-

sity unit has been defined and called Jansky (Jy), with

1 Jy = 10�26 W m�2 Hz�1. In solar radioastronomy,

and since the radio signal from the Sun is quite strong

(few orders of magnitude stronger than the brightest

non-solar radio sources in microwaves), we commonly

use solar flux units (SFU) define as 1 SFU = 104 Jy =

10�22 W m�2 Hz�1. The expression of Im radiated by ablack-body at temperature T, at the frequency m, is givenby the Planck�s law

Im ¼2hm3

c2� 1

ehm=kBT � 1� 2m2kBT

c2; ð3Þ

where we have used the Raleigh–Jeans approximation

for radio photons, i.e., hm � kBT, where h is the Planck�sconstant, kB is the Boltzmann constant and c is the

Fig. 4. Example of dynamic spectrum showing fast and slow frequency

drifting bursts. Time (in minutes) is shown on the horizontal axis, and

the observed frequency (MHz) on the vertical axis.

http://www.astro.umd.edu/white/text/activeregionimages.html





speed of light. Another quantity largely used in radioas-

tronomy is the brightness temperature Tb, which is the

equivalent temperature a black-body should have to

radiate an intensity Im given by Eq. (3). We thus have

a relation between Tb and Sm given by

Sm ¼2kBm2

c2

ZT b dX; ð4Þ

where X is the angle subtended by the emitting source.

Although the brightness temperature has been defined

for thermal emission of a black-body we can also use

it for non-thermal emitting sources. Non-thermal emis-

sion may be due to accelerated particles with energy E,

and in this case, we shall consider Tb as an effective tem-

perature Teff given by the kinetic temperature E/kB. In

the case of thermal emission, one should have Tb 6 T.We shall now distinguish between incoherent and

coherent emission mechanism. An incoherent radiation

mechanism is a process where particles behave indepen-

dently of each other. There is no phase relation between

the emitted photons, that is no coherence. In this case,

the resulting brightness temperature cannot exceed the

effective temperature Teff, that is the source temperature

for a thermal process, or the kinetic temperature for anon-thermal process. Thus for an incoherent non-

thermal emission, due to say MeV electrons as it is often

observed during solar flares, Tb is limited by the mean

particle energy, such that Tb 6 E/kB � 1010 K. During

a coherent process emitting particles can behave as a

whole, emitting photons between whose exists a phase

relation. Thus, a coherent emission is not the result of

individual particle emission, but rather a collective pro-cess for example triggered by waves when certain insta-

bilities develop in magnetized plasmas. The brightness

temperature resulting from a coherent emission mecha-

nism can thus largely exceed the mean particle energy,

and can reach few 1015 K as for some decimetric solar

radio bursts.

Due to the presence of a magnetic field an electro-

magnetic wave will propagate in two different modes:the extraordinary mode (X) and the ordinary (O) mode.

The above defined quantities will depend on the mode of

propagation. The detected radiation is said to be polar-

ized when the emission in one of the two modes domi-

nates that in the other (Kraus, 1966; Christiansen and

Hogbom, 1985; McLean and Labrum, 1985). In this

case, the degree of polarization is a measure of the inten-

sity received in one of the two modes, relative to the to-tal intensity. Since the Faraday rotation of the plane of

polarization is strong when radio waves propagate

through the solar corona, it suppresses any linear polar-

ization present at the emitting source, and thus only the

degree of circular polarization can be measured. How-

ever, see McLean and Labrum (1985) for examples of at-

tempts in detecting linearly polarized solar radio

emission.

3. Coherent radiation processes

As said earlier coherent processes may occur at the

time a magnetized plasma becomes unstable. Before

entering coherent mechanisms in more details, it is

important to remind some important parameters ofany given plasma. We will here describe few of them

and give simple expressions. For a general table of

plasma parameters relevant to solar processes, the reader

is referred to McLean and Labrum (1985, chapter 5).

3.1. Natural plasma parameters

The first parameter is the plasma frequency fpa ofplasma specie a. More often, we find in the literature,

the angular frequency xpa related to fpa by the relation

xpa = 2pfpa. The plasma frequency for electrons can be

simply written in terms of the ambient electronic plasma

density, Ne, the charge and mass of the electron

fpe ¼1

2pN ee

2

mee0

� �1=2

¼ 9� 10�3ffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiN e ½cm�3�

pMHz. ð5Þ

To understand the meaning of fpe let us assume a plasma

composed of electrons and ions on which acts an exter-

nal electric field ~E. Due to ~E there will be a charge sep-

aration which will be responsible for the creation of arestoring electric field to compensate the effect of ~E. Ifwe suddenly shut down ~E and study the motion of the

electrons, neglecting collisions between particles as well

as ions motions, we find that electrons will oscillate

around an equilibrium position with the frequency fpe.

These undamped oscillations are called plasma oscilla-

tions. In the cold plasma wave theory, these oscillations

are called Langmuir waves at the frequency x = xpe,and taking into account the electron temperature Te

(or the pressure gradient term in the equation of the

electron motion), Langmuir waves can propagate and

have a dispersion relation given by

x2 ¼ x2pe þ 3

2k2v2e ; ð6Þ

where k is the wave number and ve the mean electron

thermal velocity. It is important to note that the Lang-

muir waves are electrostatic waves, that do not transportany electromagnetic energy.

In the presence of a magnetic field another natural

plasma frequency is the gyrofrequency, fbe (for elec-

trons), defined by the frequency with which an electron

gyrate around magnetic field lines due to the Lorentz

force. fbe, often referred to Xe in the literature

(Xe = 2pfbe), can be expressed using

Xe ¼eBme

¼ 2p � 2.8� 106B½G� Hz; ð7Þ

where B is the magnetic field strength. Thus, taking into

account the magnetic field and the presence of ions, the

Fig. 5. Velocity distribution function of a thermal background plus a

fast beam, showing the velocity region where turbulence can develop.


electron plasma frequency given by Eq. (5) will be

slightly modified and will fall in the range given by the

upper hybrid frequency, xUH, and the lower hybrid fre-

quency, xLH, given by

xUH ¼ ðx2pe þ X2

eÞ1=2

;

xLH ¼ 1

x2pi þ X2

i

þ 1

XeXi

!�1=2

� ðXeXiÞ1=2.ð8Þ

In the solar corona, we generally have the relation

xUH P xpe > Xe � xLH > xpi � Xi.

The Debye length, kD is a useful parameter to study

collisions in a plasma. It defines the distance over which

a test electron from the plasma will not feel anymore the

electrostatic field from an ion. kD is the ratio of the ther-mal speed to the plasma frequency, and thus depends on

the ambient electron density Ne and the plasma temper-

ature T, through the following equation:

kD ¼ 6.9

ffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiT e ½K�

N e ½cm�3�

scm ð9Þ

being then much lower in the cool and low altitude

dense solar atmosphere compared to its value in the di-

lute and hot corona.

The plasma parameter b is a measure of the gas-to-

magnetic pressure ratio. For b � 1 like close to the solar

surface, the particles dominate the plasma dynamics and

they drag along magnetic field lines. In the outer solaratmosphere, i.e., the solar corona, we have everywhere

b � 1 implying that the plasma is confined by the mag-

netic field.

3.2. Coherent emission mechanism

The most known coherent process to produce radio

emission is the plasma emission, which is a multi-stepprocess. It first involves some kind of Langmuir turbu-

lence, generally attributed to a streaming instability.

The next step is to convert the energy available in the

plasma turbulence into fundamental electromagnetic

transverse radiation at or close to the local plasma fre-

quency (x = xpe) which can escape from the medium

and be detected. This is because Langmuir turbulence

does not propagate and can only be detected in situ.Other steps requiring the production of secondary Lang-

muir waves and the conversion into transverse electro-

magnetic emission will not be discussed here. Extended

references on these steps which involve scattering of

plasma waves on ions or on coronal density gradients,

and wave–wave interactions, can be found in Benz

(1993) and McLean and Labrum (1985).

Although the study of the system composed of abackground thermal plasma where a beam of particles

propagate is a very simplified situation, it gives already

some important details on the production of plasma tur-

bulence. The hydromagnetic treatment, that is a back-

ground plasma composed of cold particles (electrons

and ions) and a beam of cold or monoenergetic particles

(electrons with velocity Ve0), can be done graphically

(Sturrock, 1994; Kivelson and Russell, 1995). The solu-

tion of the dispersion relation given by

1 ¼x2

pi

x2þ

x2pe

ðx� kV e0Þ2

!ð10Þ

may lead to a complex solution for x of the formx = xr ± ixi. Such that any perturbation in the plasma

will varies in time as e�ixt = e�ixrtect where c = xi.

Therefore we see that c > 0 will imply an exponentially

growing perturbation. Although this simple situation

shows that under some circumstances Langmuir turbu-

lence can grow at a rate c within a plasma, it does not

tell us from where comes the energy provided to the tur-

bulence nor when the amplification will stop. To solvethis problem, it is necessary to adopt a more detailed

treatment taking into account the temperatures of the

background plasma and the beam particles, that is ki-

netic effects. In Fig. 5, we show a more realistic situation

for the velocity distribution of a background plasma

where a fast particle beam propagates. It can be shown

(Benz, 1993) that c / xrof0ovz

where xr is the real part of

the frequency given by Eq. (6). Fig. 5 shows that inthe region where of0

ovxis positive plasma turbulence with

phase velocity x/k, may suffer strong amplification.

For such a turbulence, with a phase velocity V/, we basi-

cally have more particles with velocity greater than V/

than we have particles with velocity lower than V/. So

the net effect is the growth of the turbulence, at the ex-

pense of a reduction of the beam velocity. Then, as a

function of time, the beam velocity distribution will bedestroyed, and become rather a ‘‘plateau’’. It is then a

wave–particle interaction effect, where energy is trans-

ferred from the particle beam to the waves present in

the plasma, which can explain the growth of the plasma

turbulence in a region where it exists a resonant condi-

tion between the particles and the waves. This instability

is the so-called bump-on-tail instability, which is impor-

tant to explain some kind of metric radio bursts likeType IIIs.


Instability can develop also due to the presence of an

anisotropic velocity distribution. This occurs when par-

ticles are trapped in magnetic bottles, that is a region

bounded by two sites of high magnetic field strength

Bmax. Conservation (of energy and magnetic moment)

laws show that as a particle moves toward high magneticfield regions its velocity component parallel to ~B, vi de-creases. When arriving at the point where B = Bmax,

the particle can escape the bottle if vi > 0. Otherwise

the particle will be confined in the magnetic bottle.

The mirror ratio between Bmax and the minimum field

strength within the magnetic configuration defines an

angle (pitch-angle) am such that any particle with initial

pitch-angle a ¼ ðd~v;~BÞ < am will escape from the bottle

and those particle with a ¼ ðd~v;~BÞ > am will remaintrapped. This situation is illustrated in Fig. 6 where

the pitch-angle region shown by the gray filled area is

empty of charged particles. In plasmas where Xe P xpe

the above anisotropic configuration will transfer energy

into directly escaping electromagnetic radiation,

through an instability called electron cyclotron maser

(ECM). The ECM emission (Wu and Lee, 1979) is

responsible for planetary radio emissions and some verybright solar radio spikes.

4. Incoherent radiation processes

The basis for incoherent radio emission in low density

medium like the solar corona is the emission from free

accelerated particles. A charge in rectilinear uniformmotion creates in space an electromagnetic field which

energy is static and constant in time, i.e., does not prop-

agate. When the charge is accelerated we have an addi-

tional electromagnetic field component, which is time

dependent because of the time dependent velocity. Con-

trary to the static case, the electromagnetic energy now

propagates and we say that the charge is radiating. This

is what happens for example in the emission of radiowaves when electrons oscillate in an antenna.

Fig. 6. Loss-cone velocity distribution.

The Larmor formula (Rybicki and Lightman, 1979)

expresses the power radiated by a single accelerated

charge (q, mass m, velocity v) in the direction h relative

to the acceleration vector within the solid angle dX

dPdX

¼ q2 _v2

4pc2sin2ðhÞ; ð11Þ

where _v ¼ dv=dt, and which gives after integration on X

P ¼ 2q2 _v2

3c2. ð12Þ

Note on the radiation pattern, P, that: (i) the total

power is proportional to the charge; (ii) no power is

emitted in the direction of the acceleration; (iii) maxi-

mum power is emitted in a direction perpendicular

to the acceleration; (iv) since _v is �1/m, the power is

�1/m2 and the radiation from protons will be negligiable

compared to that emitted by electrons.

4.1. Free–free bremsstrahlung

This emission is due to Coulomb collisions between

charged particles in a plasma. In Fig. 7, we show an

example of a binary collision between an electron of

velocity v and an ion of charge Zi. The effect of the col-

lision is to deflect the incoming electron by an amountwhich depends on the impact parameter b. However in

the solar coronal plasma the situation is quite different.

This is due to the very high number of particles present

in the Debye sphere (see Eq. (9)). The ratio of small-to-

large angle encounters can be approximated by kD/rcwhere rc is the impact parameter b which produces a

90� deflection. As a consequence small-angle collisions

dominate and the trajectory of an incoming electron israther determined by a multitude of small deviations.

However, in the cool and very dense low solar atmo-

sphere, the effect of large angle encounters is enhanced

since there the Debye sphere is much smaller. Then high

energy electrons can undergo large deflection and lose

most of their energy emitting high energy hard X-ray

photons. In the tenuous solar corona, the emitted pho-

tons may fall in the radio wavelength range, and thisis what we are interested in the following. In this case,

the bremsstrahlung emission from one incoming

Fig. 7. Binary Coulomb collision showing the impact parameter and

the deflection angle.


electron is calculated using the small-angle approxima-

tion (Rybicki and Lightman, 1979)

dW ðbÞdx

¼ 8

3pZ2i e

6

m2ec

3

� �1

ðbvÞ2e�2xb=v

� 8

3pZ2i e

6

m2ec3

� �1

ðbvÞ2ð13Þ

for b � v/x. Taking into account the total incoming

electron flux, 2pbdbnev, where ne is the electron density,

colliding in a plasma with an ion density ni, and integrat-

ing Eq. (13) between impact parameters bmin = 4Zie2/

pmev2 corresponding to a 90� deflection and bmax = v/x

above which the emitted power is negligiable, we get

the bremsstrahlung emission per unit time, volume and

frequency

dW ðbÞdxdV dt

¼ 16e6

3c3m2evneniZ2

i Lnbmax

bmin

� �¼ 16pe6

33=2c3m2evneniZ2

i gffðv;xÞ; ð14Þ

where gff(v,x) is the Gaunt factor (Karzas and Latter,1961) which is slightly dependent on the frequency m.For a totally ionized electron–proton plasma like the so-

lar corona (ne = ni and Zi = 1), composed of thermal

particles, we get the thermal bremsstrahlung emissivity

gm at the frequency m (x = 2pm) by integrating Eq. (14)

over v using a Maxwellian distribution function

gm ¼25pe6

3mec32p

3mekB

� �1=2

n2eT�1=2gffðv;xÞ. ð15Þ

We thus conclude that the radio emissivity from a free–

free bremsstrahlung emitting plasma is proportional to

the ambient density and inversely proportional to the

plasma temperature. We also note that as long as radia-

tive transfer of the radiation is not taken into account,

there is almost no dependence of the emission as a func-tion of the frequency.

4.2. Magnetobremsstrahlung emission

In the presence of an external magnetic field ~B, a

charged particle will spiral around field lines due to

the Lorentz force, and thus suffer continuous velocity

direction changes that will produce electromagneticradiation. In the absence of an electric field the acceler-

ation term is perpendicular to ~B, a^ = Xv^, where X is

the gyrofrequency given by Eq. (7). Using the Larmor

formula (Eq. (12)), we get the power emitted by an accel-

erated electron

P ¼ c22e2

3c3v2?e

2B2

m2ec

2; ð16Þ

where c is the Lorentz factor. In the electron reference

frame the power emitted is dipolar, and seen from a rest

frame it is beamed along the direction of ~v, also called

the forward beam. The time T, it takes for the emission

cone to pass in front the observer is 2/Xe. Because of

Doppler shift, this time s will be shorter for an observer

at rest by a factor 1 � v/c. For energetic electrons, we

have s = 2p/(c2Xe). Thus, as a function of time, thepower P(t) emitted by an electron is a succession of

pulses of width s separated by 2pc/Xe, i.e., occurring

with the frequency Xe/2pc. In the spectral domain, the

electron emission is then different depending on the Lor-

entz factor. For monoenergetic cold electrons, the emis-

sion occurs at the frequency Xe and is also called

cyclotron emission. For higher energy thermal electrons

(few eV to tens of eV) the time profile is modified andharmonics of Xe appear in the spectral domain. The

emission called gyroresonance then occurs at low har-

monics number (s = 1,2,3, . . .) of the Larmor frequency.

For mildly relativistic electrons (few tens of keV to few

hundreds of keV), P(t) presents much thinner and sepa-

rated peaks, resulting in a broadband frequency range

emission from harmonics of Xe in the range s = 10–

100. This emission is called gyrosynchrotron emission.For highly relativistic electrons, the angular pattern of

emission is beamed along the particle velocity. The emis-

sion called synchrotron emission will mainly occur along

the particle velocity at very high harmonics s � c3 of theLarmor frequency.

The treatment of magnetobremsstrahlung emission

from a collection of electrons is complex since it needs

the knowledge of the velocity and pitch-angle distribu-tions of the particles (Pacholczyk, 1979). Moreover,

the effects of the ambient medium and of an inhomoge-

neous magnetic field structure in the emitting region

should be taken into account. Complete calculations

of the emissivity for an isotropic velocity distribution

of electrons in an homogeneous magnetic field have been

performed (Ramaty, 1969; Ramaty et al., 1994). Semi-

empirical useful formulae for gm obtained within thesame conditions and valid in a restricted range of har-

monic numbers s have been estimated (Dulk and Marsh,

1982; Dulk, 1985). Inhomogeneous magnetic field mod-

els and absorption of the radiation by the ambient med-

ium have been studied (Klein and Trottet, 1984; Klein,

1987), as well as the effects of anisotropic electron

pitch-angle distribution (Lee and Gary, 2000).

Mainly two kind of particle distributions are usuallyused: Maxwellian and power-law. In the first case, the

dominant mechanism is the gyroresonance emission at

discrete and low harmonic number of Xe, which emissiv-

ity can be written as a function of the density ne, the

magnetic field strength B and its direction relative to

the observer. Thus, gyroresonance emission will be rele-

vant to study magnetic field structure above solar active

regions. In the case of a non-thermal power-law energydistribution of electrons AE�d, the gyrosynchrotron

emissivity has also a power-law spectrum which can


then be used to estimate the injected particle distribution

parameters as well as the magnetic field strength B. This

will mostly be useful to study gyrosynchrotron from

flare accelerated particles.

5. Radiative transfer

After to have reviewed how emission can be obtained

from different processes, through the emissivity term gmin W m�3 St�1 Hz�1, we will now study how this radia-

tion is transported (Kraus, 1966). The change of inten-

sity Im at the frequency m (see Eq. (3)) along a ray path

s is given by dIm = gmds. The radiation is also absorbedalong the ray path by an amount dIm = �jmImds where

jm is an absorption coefficient per unit of length. Putting

together these two expressions leads to the equation of

radiative transfer

dI mds

¼ gm � kmI m. ð17Þ

The first term is a source term while the second term is

an absorption term. The latter is used to define the opti-

cal depth sm as the integrated absorption along the line

of sight, i.e., sm ¼R s0jmds. By integrating Eq. (17) with

jm = 0 (pure emitting medium), we see that the change

in intensity is the integrated emissivity along the lineof sight. Similarly, an integration with gm = 0 (pure

absorbing medium) shows that the intensity decreases

exponentially with sm. The solution of Eq. (17) can be

written as:

Im ¼ Imð0Þe�sm þ gmjm

ð1� e�smÞ. ð18Þ

This equation means that the outgoing radiation from

an absorbing medium will be the entering radiation

Im(0) absorbed through the medium plus the radiative

balance within the medium (integrated source emissionminus absorption). At the thermodynamical equilibrium

(Eq. (17) equals 0), when the source absorbs as much as

it emits, the ratio gm/jm is given by the Planck function

(see Eq. (3)). Moreover, when dividing the former equa-

tion by 2kBm2/c2 we get

T b ¼ T b0e�sm þ T effð1� e�smÞ. ð19Þ

As already discussed in Section 2 for a source in thermo-

dynamical equilibrium Teff is the source temperature,

and for non-thermal emitting particles, Teff is to be con-

sidered as given by E/kB where E is the mean energy of

the radiating particles.

Two important regimes are found from Eq. (19): (i)

when sm � 1 the source is said to be optically thick;

(ii) when sm � 1 the source is said to be optically thin.These two cases lead, respectively, to:

T b ¼ T eff ðsm � 1Þ;T b ¼ T b0ð1� smÞ þ smT eff ðsm � 1Þ. ð20Þ

These equations will allow to compute for different emis-

sion mechanisms the variations of Tb (or Sm using Eq.

(4)) as a function of the frequency, that is to get the

radio spectrum.

For a given emission mechanism, it is common to use

the corresponding optical depth or opacity, rather thanthe emissivity, although both can be related at the ther-

modynamical equilibrium. As an example for a fully

ionized e�–p plasma we can use Eq. (15) to get the

free–free bremsstrahlung opacity (Dulk, 1985):

sff ¼9.78� 10�3 N 2e‘

m2T 3=2ð18.2þ 1.5 ln T � ln mÞ

ðT < 2� 105 KÞ; ð21Þ

sff ¼9.78� 10�3 N 2e‘

m2T 3=2ð24.5þ ln T � ln mÞ

ðT > 2� 105 KÞ; ð22Þ

where the Gaunt factor has been evaluated for two tem-perature regimes, and where ‘ is the size of the emitting

medium along the line of sight. Thus, free–free brem-

sstrahlung emission is mainly a diagnostic of the plasma

temperature and density. Note also that the radio free–

free opacity is increasing for decreasing T, i.e., very sen-

sitive to cool and dense plasmas. Because of the small

dependence of the term in brackets, the radio spectrum

(Sm as a function of m) will increase as m2 in the opticallythick regime and will be roughly constant for an opti-

cally thin source.

The opacity for the gyroresonance mechanism, sgr,can be obtained similarly as done above for the free–free

bremsstrahlung (Dulk, 1985; White and Kundu, 1997).

sgr is slightly dependent of the density and the tempera-

ture of the plasma, but strongly depends on the mag-

netic field strength and on its direction with respect tothe observer. Gyroresonance opacity is higher for the

X mode compared to the 0 mode, such that the emission

will be highly polarized. In units of the magnetic scale

height, the thickness of the gyroresonance layer is about

v/c, where v is the speed of the emitting thermal elec-

trons. Thus, the gyroresonance layer is a very thin por-

tion of the atmosphere with an almost constant

magnetic field value.For higher energy electrons with a power-law distri-

bution in energy, the spectrum of the radio flux density

decreases with the observed frequency in the optically

thin part. The peak frequency mpeak is the frequency be-

low which the radio emission is self-absorbed in the opti-

cally thick part of the spectrum. The optically thin part

of the spectrum is important since its slope is directly re-

lated to the slope of the injected distribution of electrons(Dulk, 1985). The higher the frequency the higher the

energy of the emitting electrons, and thus high fre-

quency radio observations provide a diagnostic of the

highest energy electrons accelerated during solar flares.

This is demonstrated in Fig. 8 (Fig. 1 in White and


Kundu, 1992) for different gyrosynchrotron spectra ob-

tained using different low-energy cutoff of the injected

electron distribution. One can see that millimeter emis-

sion is not affected until MeV electrons are removed,

showing that high frequency emission is largely due to

the MeV energy electrons.

Fig. 9. Variations with height of the plasma and gyromagnetic

frequencies. fs = 1 is the curve for which the free–free opacity is equal

unity (after Gary and Hurford, 2004).

6. Solar radio emissions

Before describing some selected solar radio observa-

tions, let us first have a rough idea of which process will

dominate at a given altitude in the solar atmosphere.

Fig. 9 (from Gary and Hurford, 2004) shows the varia-tions as a function of height of the plasma frequency mp,of the frequency for which sff = 1, as well as harmonics

of the gyrofrequency. At each altitude, the curve mp in-

forms us the cutoff frequency below which a wave does

not propagate. The solar plasma above the curve for

which sff = 1 is optically thin, thus, in the range 10–

200 MHz plasma emission mechanisms will dominate.

Although plasma radio emission above 200 MHz willbe partly absorbed by the ambient plasma through of

the effect of collisions, it may still be bright enough to

be detected. The reason is that plasma radio waves are

naturally extremely bright and can occur at harmonics

of the plasma frequency. The highest frequency at which

plasma radio emission may have been detected in the so-

lar corona is about 8 GHz. Thus, for frequencies be-

tween 200 MHz and �1 GHz, both mechanisms,plasma and bremsstrahlung can coexist. Between 1 and

3 GHz the hot and dense plasma from the active regions

is optically thick for free–free bremsstrahlung emission

which is then the dominant mechanism. Above about

3 GHz, the altitude for which the ambient plasma has

Fig. 8. A plot of non-thermal gyrosynchrotron flux spectra as a

function of frequency for different values of the low-energy cutoff. The

spectral index is 4, the magnetic field 300 G, the high-energy cutoff is

10 MeV and the angle between the line of sight and the magnetic field

is 45�. The curves are labelled according to the low-energy cutoff (after

White and Kundu, 1992).

a free–free opacity equals to 1 is lower than that corre-

sponding to low harmonics of the gyrofrequency. Con-

sequently, gyroresonance mechanism will prevail in

high magnetic field regions.

6.1. Non-flare radio emissions

The oldest observed solar radio emission at metricwavelengths is called a noise storm and has been re-

viewed by Elgaroy (1977). It is composed of short dura-

tion (61 s) metric radio enhancements called Type I

bursts, superimposed on a broad frequency band contin-

uum (McLean and Labrum, 1985). Noise storms are

spatially associated with active centers, but are not flare

related (Le Squeren, 1963), and they can last for several

consecutive days. Their importance lies on the fact thatthey are evidences for long lasting acceleration of elec-

trons up to suprathermal energies, which produce radio

emission through a collective process (Spicer et al., 1982;

Wentzel, 1986). Noise storms are triggered in a similar

way as flares (Raulin et al., 1991; Raulin and Klein,

1994), the emitted power being much less than that mea-

sured during flares, although the total energy budget is

comparable to that of a small flare. We finally note thatType I bursts may be the coronal signature of nanofl-

ares, and their peak flux density distribution suggests

that they may contribute to the heating of an active

coronal region (Mercier and Trottet, 1997). The slope

of the distribution, for small events, is much steeper than

that found at other wavelengths, indicating the possibil-

ity that Type I bursts participate to the heating of the so-

lar atmosphere.The quiet sun also emits at radio wavelengths

through the free–free bremsstrahlung mechanism. How-

ever the emission will greatly depend on the atmospheric

layer because of the dependence of the free–free opacity

sff (see Eqs. (21) and (22)), on the plasma density Ne, and

the local temperature T. Using a typical temperature of

Fig. 10. Optical (left) and radio (right) emissions above a sunspot.

Note the ‘‘horse-shoe’’ like structure of the radio source (Courtesy of

S.M. White, after http://www.astro.umd.edu/white/text/activeregio

nimages.html).


106 K and density of 109 cm�3 we get sff � 1 in the

whole microwave range indicating that the quiet corona

will be optically thin at these frequencies. Through the

transition region to the chromosphere, Ne increases by

a factor of �100 and T decreases by about the same fac-

tor, implying a change of seven orders of magnitude insff resulting in an optically thick chromosphere. Zirin

et al. (1991) have effectively shown that the free–free

emission from the quiet sun center between 1.4 and

18 GHz was composed of the radiation of a thick upper

chromosphere at �11,000 K, observed through a thin

corona at 106 K and scale height of 5 · 109 cm. At meter

wavelengths the situation is somewhat different since the

corona will be optically thick at some height showingbrightness contrasts with extended regions called coro-

nal holes. These regions are cool and low density areas

in the corona, and their time variability is important

to characterize the high-speed solar wind. At much

higher frequencies (millimeter and submillimeter wave-

lengths), the radio emission is sensitive to the cool and

dense chromospheric material which provide enough

free–free opacity, while the corona remains totally opti-cally thin. Vernazza et al. (1981) tested successfully their

chromospheric atmosphere model deduced from EUV

observations, against 100 GHz radio measurements of

the quiet Sun which originate at layers where the tem-

perature was �6600 K. However, the same model atmo-

sphere seems to disagree with optically thick microwave

observations (Zirin et al., 1991; Bastian et al., 1996). On

the other hand, the latter radio observations are well ex-plained by chromospheric models based on CO line

measurements (Avrett, 1995), suggesting that a chromo-

spheric model which could explain, high frequency

radio, microwaves, EUV and line observations is still

missing.

Solar active regions are high density and high mag-

netic field regions located above sunspots. Radio emis-

sion will occur due to the two following mechanisms:gyroresonance and free–free bremsstrahlung. Which

mechanism will dominate depends on the density and

the magnetic field strength in the source, and the fre-

quency of observation. The former mechanism is

strongly dependent on the magnetic field strength, but

less on the plasma parameters, while free–free emission

mainly depends on the plasma temperature and density.

Fig. 9 has shown that below a frequency of a few GHzthe local plasma is optically thick to free–free emission.

Above few GHz and in the presence of strong magnetic

fields, low coronal regions of constant magnetic field

provide high gyroresonance opacity. As already men-

tioned sgr is only significant in thin coronal layers (Dulk,

1985; White and Kundu, 1997) where the observed fre-

quency is equals to low (s = 1, 2, 3, 4) harmonics of

the gyrofrequency. Therefore, providing an uniqueway of measuring magnetic field strength. The structure

of the gyroresonance emission source depends also

strongly of the angle h, between the magnetic field and

the line of sight, thus, explaining horse-shoe emission

shapes like the one shown in Fig. 10. Sunspots micro-

wave radio emission compared to model calculations

(Gopalswamy et al., 1996; Nindos et al., 1996) have con-

firmed that gyroresonance is the main mechanism atwork.

6.2. Solar flare radio emissions

Solar flares produce copious amount of coherent

radio waves, which have been classified for more than

40 years into different classes (see Kundu, 1965; McLean

and Labrum, 1985; Benz, 1993, for reviews). The mainobservational parameters for this classification are those

that can clearly be identified on dynamical spectra re-

cords (see Fig. 4), i.e., bandwidth, frequency drift rate

and duration of the emission. The different types of

coherent bursts in the decimetric/metric domains will

not be discussed here and the reader is referred to the

above reviews. Instead we rather describe some charac-

teristics of Type III solar bursts in the following para-graph, while properties of Type I solar bursts have

already been discussed in the previous section. We also

mention that there is no clearly accepted emission mech-

anism for any of the coherent solar radio bursts, and the

non-linear nature of the coherent process itself makes it

very difficult to conclude on quantitative estimations

like electron number and spectrum which represent cru-

cial informations to test particle acceleration models(Miller et al., 1997).

Type IIIs are among the more studied solar coherent

bursts, and generally occur at the beginning of the

impulsive phase of a solar flare. As seen in dynamic

spectra, Type IIIs are fast drifting bursts (see Fig. 4)

from high-to-low frequencies, i.e., from high critical

plasma frequency levels to low critical plasma fre-

quency levels. Reverse bursts are those which will showfast drifts from low-to-high frequencies. Type IIIs are



Fig. 11. Time profile of a small flare at microwave (17 GHz; dashed),

hard X-rays (histogram) and at millimeter-waves (86 GHz, BIMA,

thick) (after Raulin et al., 1999).


thought to be produced by plasma waves excited by

electron beams propagating at different levels of the so-

lar atmosphere. Since electron beams propagate along

magnetic field lines, Type IIIs are used as tracers of

magnetic field structures. For example radio Type IIIs

are observed during the rearrangement of the magneticfield structure within active centers resulting in colli-

mated plasma outflows or ‘‘jets’’. In that case the elec-

tron beams, as imaged using multi-frequency radio

observations, nicely align with the jet material (Kundu

et al., 1995; Raulin et al., 1996). The jet is observed in

soft X-rays, i.e., a diagnostic of the ambient electron

density which can be compared to that deduced from

the observed radio frequencies, thus, providing infor-mations on the emission mechanism (Raulin et al.,

1996). Similarly, different spacecrafts have been used

to reconstruct the trajectory of a Type III which es-

capes into the interplanetary medium (Reiner et al.,

1998). The electron beam path is following a Parker

spiral (i.e., interplanetary magnetic field structure),

and the speed of the exciters electrons was measured

at about �100,000 km/s.Bi-directional Type IIIs are identified on dynamic

spectra as positive and negative (or reverse) fast fre-

quency drift bursts starting at a given frequency range.

Such analysis has been reported by Aschwanden et al.

(1993), which allows to conclude that the acceleration

region was located at a plasma level of 300–500 MHz,

corresponding to an ambient density of 3 · 108–

3 · 109 cm�3. More recently, other informations on theacceleration and energy release sites were presented by

Paesold et al. (2001). In this work, double Type III

sources were found to originate from the same metric

spike source, suggesting that the fast spikes were closely

related to the acceleration region.

As mentioned earlier, part of the energy dissipated

during solar flares is used to accelerate particles to

high energies. As part of these, electrons propagatein the solar atmosphere emitting radio waves through

the gyrosynchrotron mechanism when passing by high

magnetic field strength regions, and eventually collide

in the dense lower atmosphere where they are instan-

taneously stopped and emit high energy hard X-ray

photons through the non-thermal bremsstrahlung pro-

cess. Both of these mechanism thus provide direct

diagnostics of the energetic electrons. However, andwhen compared to non-thermal bremsstrahlung, the

gyrosynchrotron process is quite sensitive to the pres-

ence of even only a few numbers of energetic elec-

trons, as long as they are spiralling in strong

magnetic fields. This property allowed to detect gyro-

synchrotron emission from very simple small flares

(Kundu et al., 1994; Raulin et al., 1999), with flux

density levels as low as 1.5 SFU, showing that flaresof all sizes are capable of accelerating electrons to

MeV energies. Since both processes depend on the

particle energy, but only the gyrosynchrotron emission

is sensitive to the magnetic field, their comparison can

bring clues on the transport and pitch-angle distribu-

tion of the emitting electrons, as well as on the mag-

netic field strength. The transport of microwave

emitting electrons has been studied by Lee and Gary(2000) and tested against spatially resolved radio and

hard X-ray observations. As expected it is shown that

the pitch-angle distribution of the injected electrons

precludes on the temporal evolution of the radio emis-

sion, and a rather beamed (a 6 30�; see Section 3.1)

injection distribution is needed. On the other hand,

an example of efficient trapping is illustrated in Fig.

11 (Raulin et al., 1999), which shows that the impul-sive phase of a small flare is composed of a fast peak

seen in microwaves, millimeter-waves and hard X-ray,

and a few minutes duration emission intense in radio

wavelengths but with no counterpart in hard X-rays.

The electron energy spectrum at the time of the first

rapid peak was harder at energies of �1 MeV (i.e.,

microwave and mm-w emitting electrons), than at

6 few 100 s keV (hard X-ray emitting electrons). Thisresult has been observed during large flares (Trottet

et al., 1998), and also during smaller events (e.g.,

Kundu et al., 1994). However, in the absence of imag-

ing informations, it is not possible to conclude

whether the observed broken spectra are the result

of a single acceleration process, or if the low and high

energy parts of the spectra are due to independent

injections, for example localized at different places inthe solar atmosphere.

Fig. 12. Radio spectrum at different peaks (P1 and P4) during the

impulsive phase of the 2003, November 4 large solar event. Spectrum

for fast submillimeter structures during P1 is also shown (after

Kaufmann et al., 2004).


7. Solar radio emission at very high frequencies

The electromagnetic spectrum above 100 GHz has

been mostly unexplored until the last years. Early micro-

wave solar flare observations have shown radio flux

spectra increasing with frequency, suggesting a spectralmaximum above 30 GHz (Hachenberg and Wallis,

1961), which could represents the superposition of differ-

ent synchrotron spectra or the emission from ultra-

relativistic electrons. At higher frequencies, there are

numerous examples of complex synchrotron spectra ob-

tained during both large and small solar bursts showing

flat or increasing fluxes with the observed frequency (see

Fig. 2 of Kaufmann, 1996, for a summary). These exam-ples suggest a spectral peak located above the highest

observed frequency of �100 GHz. Few solar observa-

tions were made at even higher frequencies: at

250 GHz (Clark and Park, 1968) reported one minute

temporal fluctuations of the order of �100 K above ac-

tive regions, and at 850 and 12,000 GHz 10–50 K varia-

tions were detected by Hudson (1975).

A new solar submillimeter facility was conceived to fillthe gap of observations above 100 GHz (Kaufmann,

1996). The solar dedicated Solar Sub-millimeter Tele-

scope (SST) was installed in the argentinean andes in

1999, May (Kaufmann et al., 2001a). The SST has four

212 GHz and two 405 GHz radiometers located at the

focal plane of a 1.5 m diameter reflector, recording with

5 ms time resolution. The beam disposition allows to

estimate the position of the centroid emission for a smallsource, using the multiple-beam technique (Gimenez de

Castro et al., 1999). The SST was operating during short

solar observing campaigns till 2001, April when daily

observations began. The main results obtained for the

solar events analyzed so far are: (i) the existence of fast

subsecond pulses (Kaufmann et al., 2001b, 2002; Raulin

et al., 2003; Makhmutov et al., 2003), superimposed on a

slower bulk component emission which is the high-frequency tail of a microwave gyrosynchrotron spectrum

(Trottet et al., 2002); (ii) when both fast and bulk compo-

nents are observed, the rapid pulses� occurrences rate andamplitudes time profiles follow the time history of the

main bulk emission at submillimeter-waves, and at hard

X-rays and c-rays energies (Kaufmann et al., 2002; Rau-

lin et al., 2003); (iii) spectra of individual fast structures

are increasing, or flat, between 212 and 405 GHz (Kauf-mann et al., 2001b; Raulin et al., 2003). More recently,

Kaufmann et al. (2004) presented submillimeter, milli-

meter and microwave observations obtained during the

large solar flare of 2003 November, 4. The properties

of the radio emission agree with the above mentioned

properties (ii) and (iii). However, the main bulk emission

at submillimeter-waves was clearly not the high fre-

quency part of the optically thin gyrosynchrotron micro-wave/millimeter spectrum, as shown in Fig. 12. Indeed

the bulk flux density has been found to increase between

212 and 405 GHz. Suggesting the detection of a new so-

lar flare emission component radiating in the terahertz

range only. These observations imply strong constraints

on the acceleration mechanisms, on the radiation pro-

cesses and the source parameters. In a non-thermal syn-chrotron interpretation, the high turnover frequencies

would imply ultra-relativistic emitting electrons, high

magnetic field values, and extreme conditions in the emit-

ting region, like ultra compact sources and/or high ambi-

ent densities, as to produce turnover frequencies in the

submillimeter wavelengths range or higher, as during

the 2003, November, 4 solar event. A thermal interpreta-

tion would require high density and cool plasma regions,like the chromosphere, to produce turnover frequencies

at P400 GHz. But at the same time the high flux densi-

ties observed (e.g., during the 2003, November, 4 event)

need to be explained by a rather large emitting source,

and this was not the observed.

Radio observations obtained in a large frequency

range from 1 to 405 GHz has also been used to estimate

the number of accelerated electrons and magnetic fieldstrength in the flaring region. Raulin et al. (2004) report

the study of the 2001, August 25 X5.6 solar flare which

produced over 100,000 SFU at 89.4 GHz. The interpre-

tation of the radio spectra during the impulsive phase,

as due to synchrotron emission from�20 MeV electrons,

requires that 5–6 · 1036 electrons with energies P20 keV

are accelerated each second and radiate in a 1000–

1100 G region. The self-attenuation of the 89.4 GHzemission during the most intense part of the event is com-

patible with a high density of non-thermal electrons

(�1011 cm�3 electrons above 20 keV). A suggestion is

that non-thermal electrons were accelerated in the low

solar atmosphere, scenario which also agrees with the


fact that the most intense part of the event was triggered

along with the interaction of very compact low-lying

magnetic loop systems. Such high non-thermal electron

numbers were also recently deduced using radio observa-

tions up to 80 GHz and hard X-ray data from the RHES-

SI detectors (White et al., 2003). In any case, the highnumbers deduced from these observations indicate that

the acceleration mechanism should be very efficient.

Solar flare observations have also been performed at

210, 230 and 345 GHz using the KOSMA radiotelescope

(Luthi et al., 2004a,b), during two X class flares. The

more relevant result of these studies is the finding of

the source size at 212 GHz of about �1000 during the

impulsive phase of one of the events. During the gradualphase of few tens of minutes duration, the radio source

size was as big as 1 0.

8. Space weather related solar radio emissions

During the last two decades it became clear that solar

conditions greatly affects the Earth and its environment.The solar influence acts through particles, radiation and

mass ejections, and their time variabilities define the

space weather conditions. Extreme space weather condi-

tions can in turn have practical consequences for life on

Earth and in space. Radio observations indices such as,

245 MHz burst and noise storm fluxes, 10.7 cm flux, are

used in Space monitoring by the Space Environment

Center (NOAA) to produce daily alerts. Occurrence ofType II bursts (see McLean and Labrum, 1985, for a re-

view), which are characteristics of radio emissions asso-

ciated with propagating disturbances in the corona and

interplanetary medium, are also used for this purpose.

Coronal mass ejections (CMEs) in particular have a

great influence on the interplanetary medium, producing

shocks and possible geomagnetic storms at the time of

their impact with the Earth magnetic field. Along withthe destabilization of large scale magnetic structures that

form a CME, prominence eruptions are often detected

(Gopalswamy et al., 2003) and sometimes well associ-

ated with the bright core of the CME. The study of plas-

ma ejections in radio waves allows to observe the

initiation of the ejection close to the solar surface, con-

trary to what can be done using coronographic observa-

tions. Also, white-light emission is only sensitive to thedensity integrated along the line of sight, while radio

emission can detect possible non-thermal particles asso-

ciated with the ejection.

Prominence eruptions have been extensively studied

at radio wavelengths, in particular at 17 GHz (see

Gopalswamy et al., 2003 for a statistical study) by the

Nobeyama Radioheliograph (Nishio et al., 1994). The

emission mechanism is optically thick free–free brem-sstrahlung (see Eqs. (21) and (22)) from the cool and

dense prominence plasma, which generally appears in

depression onto the solar disk and bright over the limb.

Then the erupting material can be tracked in the corona

up to distances over 3 solar radius. More recently, Kun-

du et al. (2004) have shown an example where a filament

eruption starts far away (�0.5 Rx) from the site of the

associated flare, implying that if both phenomena arephysically related, the destabilization of a large scale

coronal magnetic field is required.

Secondary radio products of CMEs have also been

observed in the form of Type II shocks, Type IV (see

McLean and Labrum, 1985) and storm continua (see

Pick, 1999, for a review). Bastian et al. (2001) report

the imaging of a CME at decimeter-meter wavelengths

and in white-light using the Nancay Radioheliograph(Kerdraon and Delouis, 1996) and the LASCO experi-

ment (Brueckner et al., 1995), respectively. The authors

show that the CME expanding arches are illuminated by

�0.5–5 MeV electrons emitting synchrotron radiation in

a �0.1 G, and that these particles could have been pro-

duced during the associated flare. Such observations

provide in principle independent diagnostics of the

CME magnetic field as well as of the ambient electrondensity. The signatures of the interaction between two

CMEs can also be detected. Gopalswamy et al. (2001)

have studied the trajectories in the interplanetary space

of subsequently ejected CMEs and noted that at about

the time the trajectory paths intersect each other, radio

enhancements are observed. The radio emission is inter-

preted as the result of the collision between the two

CMEs. The authors comment on the relevance of theirobservations in the context of Space Weather and for

predicting the arrival of Earth directed perturbations.

Recent observational suggestions have been reported

which might bring new informations on the processes at

the origin of these energetic phenomena that are CMEs.

Kaufmann et al. (2003) presented a study on the associ-

ation between the launch time of CMEs detected by

LASCO and the occurrence rate of fast submillimeterpulses observed at 212 and 405 GHz by the SST. Few

of the CME examples presented were not associated

with a flare on the solar disk. By extrapolating the

CME positions back to the solar surface, the onset of

the ejections were nearly coincident with a significant in-

crease of the submillimeter pulsed activity. Although the

characteristics of the fast submillimeter bursts have al-

ready been reported (Makhmutov et al., 2003; Raulinet al., 2003), their origin is still unknown. They might

represent multiple small scale instabilities, and an early

signature of the launch and acceleration of large masses

of coronal plasma.

9. Conclusions

We have presented a tutorial review on solar radio

emissions where we have defined some radio astronomy


fundamentals and introduced basic principles about

coherent and incoherent processes. Selected examples

of radio observations were chosen and discussed in

terms of local physical solar plasma parameters like elec-

tron density and temperature. We paid particular atten-

tion to explosive high frequency radio observations inthe millimeter/submillimeter wavelength range, because

this frequency range is sensitive to the highest energy

synchrotron producing electrons, and since a good

description of the optically thin part of the emitted spec-

trum is needed to find out the properties of the emitting

particles. However, the situation may not be as simple as

this classical scheme in the view of the explosive submil-

limeter emissions. In any case, further similar diagnos-tics are needed and some well-known mechanism may

need to be revisited. At lower frequencies, a good spa-

tial, temporal and spectral description of the coronal re-

gion where it seems that the acceleration processes occur

is also needed.

The new proposed Frequency Agile Solar Telescope

array (FASR) will answer part of the questions raised

above (Bastian, 2003). This new instrument will be com-posed of about �100 antennas, and will detect solar

radiation in the 0.3–30 GHz frequency range, with high

angular, temporal and spectral resolutions. FASR will

essentially address solar flare physics and active region

magnetic structure questions. At low frequencies, and

combined with other existing radio facilities, FASR will

be very suitable to study the initiation of CMEs, and de-

scribe them as a whole phenomena, thus, contributing inan important way in our understanding of Space

Weather and near-Earth conditions.

Acknowledgments

This research work was supported at Universidade

Presbiteriana Mackenzie by the MACKPESQUISAprogram and by Brazilian agency CNPq (contract

300782/96-9). The SST is supported by the Brazilian

agency FAPESP (contract 99/06126-7), and the Argen-

tinean agency CONICET. The support and help re-

ceived by the engineers and technicians from CASLEO

at El Leoncito Observatory was essential for the success

of the observations obtained with the SST. A.A. Pacini

thanks the members of the organizing committee of thefirst Advanced School on Space Environment (ASSE

2004) for the support she received.

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