PHYSICAL REVIEW D 74, 052003 (2006)

Did the 28 October 2003 solar flare accelerate protons to * 20 GeV? A study of the subsequent
Forbush decrease with the GRAPES-3 tracking muon telescope
T. Nonaka, Y. Hayashi, N. Ito, S. Kawakami, T. Matsuyama, A. Oshima, H. Tanaka, and T. Yoshikoshi
Graduate School of Science, Osaka City University, Osaka 558-8585, JapanS. K. Gupta,* A. Jain, S. Karthikeyan, P. K. Mohanty, S. D. Morris, B. S. Rao, K. C. Ravindran, K. Sivaprasad,
B. V. Sreekantan, S. C. Tonwar, and K. Viswanathan
Tata Institute of Fundamental Research, Homi Bhabha Road, Mumbai 400 005, India

H. Kojima
Nagoya Women’s University, Nagoya 467-8610, Japan(Received 14 February 2005; revised manuscript received 10 April 2006; published 18 September 2006)
Solar flares accelerate charged particles through a variety of mechanisms, which may be constrained
through observations at high energies ( > 10 GeV). We report here a search for direct emission of protons
of energy * 20 GeV in association with an X17 class solar flare that occurred on 28 October 2003, using alarge area tracking muon telescope of the GRAPES-3 experiment at Ooty. Some features of the telescope,
including its novel capability of high sensitivity search for the directional enhancement of the solar
protons are also described. A 99% C.L. upper limit on the flux of protons due to the solar flare has beenplaced at 1:4  106 cm2 s1 sr1 . A separate upper limit on the narrow solid angle flux of protons at
4  106 cm2 s1 sr1 is also placed. Solar flares are also associated with coronal mass ejections, which
propagate through the interplanetary space producing geomagnetic storms and Forbush decrease (Fd)events, upon their arrival at the Earth. New information on the structure and time evolution of the large Fd
observed on 29 October 2003 by GRAPES-3 is presented. The onset of Fd in nine different solid angle
bins (  0:3 sr) shows a remarkably similar behavior, with an evolution on a time scale of 1 h. A powerlaw dependence of the magnitude of the Fd on the cutoff rigidity has been derived, using the data fromtracking muon telescope, over a narrow range of cutoff rigidity 14.3–24.0 GV, which shows a spectral
slope ‘‘  0:53  0:04,’’ in agreement with earlier measurements.
DOI: 10.1103/PhysRevD.74.052003

PACS numbers: 96.50.Vg, 96.50.S, 96.60.qe

I. INTRODUCTION
The Sun is the closest astrophysical source accelerating
particles to GeV and possibly higher energies, yet there isvery limited understanding of, and little consensus on, thesites and processes leading to the production and acceleration of energetic particles associated with the solar flare
activity. Here, we report on the search for the emission of
protons of * 20 GeV during the active phase of the large
X17 class solar flare that started at 09:51 UTC (coordinateduniversal time) on 28 October 2003, using a large area
tracking muon telescope of the GRAPES-3 experiment.
Protons were accelerated during the solar flare on
28 October 2003 around 11:10 UTC nearly 80 min after
the detection of this flare in the x rays. This solar flare had
occurred in the active sunspot region 10 486 in the east of
the central meridian of the Sun [1].
An intriguing feature of the solar flare of 28 October2003 has been that the earliest arriving particles wererecorded by the neutron monitors (NMs) viewing the antisunward direction. Also the enhancement in the nucleon
flux continued for a long duration of time (  20 h) until
the fast moving (mean velocity  2200 km=s) coronalmass ejection (CME) shock front arrived at the Earth.
*Electronic address: gupta@grapes.tifr.res.in

1550-7998= 2006=74(5)=052003(15)Two distinct components, namely, a prompt and a delayed
one, were observed by the global network of Earth-based
NMs [2]. Detailed modeling of charged particle propagation in the interplanetary medium permeated with magnetic fields to explain the delayed and the prompt
components of the solar particles, observed by the worldwide network of a large number of NMs, has been attempted by Miroshnichenko and co-workers [3]. Theseauthors have exploited the delays in the onset of the rateincreases seen in the NMs to estimate that the path lengthof the relativistic protons from the Sun to the Earth to be2:2 AU. This model implies that the acceleration of the
protons in the flaring region is followed by their transport
along a stretched magnetic field loop (  2:2 AU) to the
Earth. During this period the Earth was located at theinterface of the interplanetary CME and a corotating
stream from the Sun. The direction of the loop from theantisunward direction was used to explain the direction of
the radiation and the anisotropy seen in the NM data.
The onset of the Forbush decrease (Fd) started in thewake of a fast CME [4], on 29 October 2003 nearly 20 hafter the start of the X17 class solar flare at 09:51 UTC on
28 October 2003. Concurrent observations of sub-GeV
solar neutrons using several NMs and of solar  raysfrom INTEGRAL satellite have been made with good
statistical significance. The correlation in the time profiles

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T. NONAKA et al.

PHYSICAL REVIEW D 74, 052003 (2006)of the solar neutrons and the  rays is used to conclude thatthe solar neutrons and the  rays were produced at the sametime [5].
The available data obtained so far have shown that thesolar flares and CMEs are a complex combination of avariety of physical processes involving the reconnection of
the solar magnetic field lines due to certain instabilities
occurring in the chromospheric and coronal regions surrounding the Sun. In some cases, the solar flares areobserved in association with the CMEs. However, nosuch association between these two phenomena is seen ina number of other instances. The relationship between theemission mechanisms responsible for the solar flares and
CMEs is an area of active research and much debate (seeRef. [6] and references therein). Solar flares are among themost energetic events to occur in the solar system, releasing vast amounts of energy (  1025 J) on relatively shorttime scales (  102 s). Solar flares give rise to the emissionof a variety of radiation ranging from radio, microwave,
optical, soft and hard x rays, and  rays of up to several
MeV including the  ray lines due to nuclear deexcitation.
Energetic solar particles contain electrons and ions which
are accelerated in the solar flares [7].
Solar flares are classified into two broad categories,
namely, ‘‘gradual’’ and ‘‘impulsive’’ flares. The gradualflares occur higher up in the corona and generally last for alonger duration of time and are also spread over a larger
geographical region on the Sun. The gradual flares are also
associated with coronal mass ejections and are accompanied with soft and hard x-ray emission. The ions accelerated in the gradual events display a composition which israther similar to that of the coronal abundances. The impulsive flares on the other hand are believed to emanatefrom the regions closer to the Sun in the chromosphere and/or lower corona; the emission takes place for a shorter
duration of time and over a relatively compact geographical region. The composition of the energetic ions in theimpulsive events displays a very significant enhancementof the heavier nuclei as compared to the coronal abundances by an order of magnitude or more. In the case ofhelium, the ratio 3 He=4 He is observed to be 1 as compared to the coronal value of 103 [7,8].
Based on the observations, it appears that the ions in thegradual events are accelerated by a CME driven shock
front. The shock front itself may be formed as a result of
the magnetic reconnection, induced due to some instability
in the coronal region. In the case of the impulsive events,
several processes have been proposed for the modeling of
the turbulent acceleration of particles in the flare region.
Stochastic acceleration involving interaction of the particles with plasma waves gives a reasonable fit to some of thefeatures of the energy spectra and the abundance ratios forseveral ions in the impulsive solar flares [7,9]. However,some of the recent observations have shown that thiscanonical scenario of different acceleration mechanisms

and characteristics for the impulsive and gradual flares is
not always valid. It has been argued that the mechanism(s)responsible for impulsive events may also accelerate ionsin some gradual events leading to a heavy ion rich impulsive composition in those cases [10]. Alternatively in someinstances, shock acceleration could play a significant rolein the acceleration of the particles in impulsive events. Thishypothesis has been borne out by the observation of a CMEin association with an impulsive flare on 1 May 2000 [11].
While the primary particles generated in a solar flare arecharged, and are accelerated in the solar convection zones
which serve to transfer energy to the interaction regions in
the chromosphere, the  ray and neutron emission provides
information on the interaction between solar energetic
particles and the ambient medium around the Sun. For
example, the temporal profiles of the  ray flux at different
energies provide a wealth of data on the interaction ofvarious particles and their entrapment near the Sun. The
continuum  ray emission is produced due to the bremsstrahlung from high energy electrons. The bombardment of
the chromosphere by high energy ions leads to the excitation of the ambient matter resulting in the emission of  raylines in the MeV region due to the nuclear deexcitation in
the chromosphere [12]. On the other hand, solar neutrons
provide crucial information on the interaction of energetic
ions, as the neutrons escaping into the interplanetary medium can be observed as decay protons or electrons by thedetectors in space, and in rare cases, through the recoil
protons in the ground-based NMs. However, the NMs
mostly observe secondary neutrons produced by the interaction of primary protons with air nuclei.
In terms of the requirement for the energy budget, theconversion of the magnetic potential energy into otherforms, during the process of magnetic reconnection [13]provides a viable and attractive possibility to accelerateparticles to several tens of GeV in a relatively short time.
Therefore, there is a great deal of interest in exploring theenergy spectrum of the particles emitted during the solar
flares, over a broad range of energy. This becomes especially important in view of the fact that the NMs operating
at high rigidities have only rarely recorded high energy
events above 10 GeV. As in the other branches of astrophysics, a significant advance in the understanding of thephenomenon of the solar flares is unlikely to occur in theabsence of multiwavelength observations with a variety ofground- and space-based instruments. This approach has
been successfully employed in the case of the large solar
flare that had occurred on 29 September 1989 [12].
While the satellite-borne detectors with high sensitivity
such as the instruments aboard the Advanced Composition
Explorer [14] provide excellent measurements at sub-GeV
energies, the energy range spanning from a few GeV totens of GeV, can only be explored through the studies, with
the network of NMs and muon telescopes located at various sites around the world. The highest energy of the052003-2

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PHYSICAL REVIEW D 74, 052003 (2006)

particles accelerated during the solar flares is one of thedecisive parameters that can be used to differentiate among
various possible models of particle acceleration in the Sun.
Observations of flare particles show spectra with cutoff atenergies beyond a few GeV which have to be understoodon the basis of the models of particle acceleration.
Observations carried out using the detectors with highsensitivity opens up the possibility of search for particle
fluxes at high energies near the cutoff and/or the shape ofthe energy spectrum near such cutoffs. Such measurements
may provide new constraints on models of particle acceleration in solar flares. However, the galactic cosmic rays
constitute the dominant background for the detection of thehigh energy (above tens of GeV) solar particles. The
detection of the secondary muons (produced in the atmosphere by energetic solar particles) with the surface and theunderground muon telescopes with good directional capability is an important tool to probe the high energy end
point of the energetic solar particle spectrum. It has been
pointed out that muon detectors with relatively good angular resolution may be able to observe the energetic solar
particles at 100 GeV.
In the past, there have been several reports of the detection of the bursts of muons which were nearly coincident intime with the associated solar flare activity in x rays,
although with somewhat limited statistical significance(see [12] and the references therein). One of the significant
detection of a muon burst was by the Baksan group which
occurred on 29 September 1989 at an energy of *200 GeV, indicating possible acceleration of the solar
protons to several hundreds of GeV. However, the directionof this muon burst did not coincide with the direction of theSun. Also the enhanced muon flux was observed after adelay of a couple of hours after the impulsive phase of theassociated solar flare [15]. The delayed arrival of the highenergy solar particles opens up the possibility, in at least
some of the solar flares, of the acceleration of high energy
protons and ions after a significant delay, following theinitial impulsive phase. The AGASA group had reported
the detection of an enhanced flux of prompt muons during
the solar flare of 4 June 1991, which they claimed was dueto the arrival of the solar neutrons of energy  10 GeV[16]. The L3 group, using their precision tracking muon
detectors at CERN, have reported the detection of a significant excess of muons in association with the solar flareof 14 July 2000 [17].
The occurrence of the solar flares and the CMEs result ina sudden injection of high energy plasma with frozen-in
magnetic field lines into the interplanetary space. Thisphenomenon leads to the formation of a shock front in
the interplanetary space surrounding the Sun. The electromagnetic field present in the plasma acts as a potential
barrier for the incoming cosmic radiation, when the shock
front passes in the vicinity of the Earth. As a consequence,
the observed rate of the secondary particles (neutrons andmuons etc. produced by the primary cosmic rays) in theatmosphere of the Earth registers a sudden drop, which istermed ‘‘Forbush decrease’’ [18]. The passage of the Earththrough the shock front results in the sudden commencement of a geomagnetic storm. Several variants of the shockmodels involving ordered and/or turbulent magnetic fields
have been proposed to explain the phenomenon of Fd [19].
The Fds are characterized by a rapid decrease ( 
5%–15%) in the intensities of the observed muon andneutron fluxes at the Earth with an onset time, ranging
from several hours to about a day. This is followed by arelatively slower recovery time, ranging from about a week
to several months. The magnitude of a Fd is defined as thefractional decrease in the intensity of the muon/neutron
flux from its quiescent value. The magnitude of the Fds isobserved to decrease with increasing rigidity. The Fds also
display some correlation with the 11-year solar cycle. The
occurrence of the Fds is generally accompanied by theobservation of large anisotropies in the flux of the cosmic
radiation [19]. The phenomenon responsible for a Fd
evolves as it propagates farther into the heliosphere. Thisevolution results in a gradual increase in the onset and therecovery times with the heliocentric distance. There havebeen suggestions that the cosmic ray modulation with the
11-year solar cycle could be a result of the cumulative
effects of the Fds occurring over a long period of time [20].
The ground-based NMs and space-borne detectors ingeneral have a very poor angular resolution. Therefore,
the information on the directional dependence of the Fds is
usually obtained from the worldwide network of groundbased NMs located all over the Earth, with the help of
appropriate modeling techniques [21,22]. However, due tothe variation in the cutoff rigidity with the geomagnetic
latitude of the site, various NMs respond differently to agiven Fd event, thus making the interpretation of theobserved data somewhat complicated and model dependent. However, a single monitor with a good angularresolution may be able to provide simultaneous information on the directional dependence of the observed Fd. In
addition, the detectors located near the geomagnetic equator have rigidity cutoff, which varies rapidly with thearrival direction, thus permitting the measurement of therigidity dependence of the detected Fd from a single
instrument.
The arrival of the CME on 29 October 2003, 20 h after
the start of the X17 class solar flare at 09:51 UTC on
28 October 2003, started the onset of a large Fd. As
described later in Sec. V, new information on this Fd has
been provided by the GRAPES-3 experiment. The magnitude of this Fd as measured by the GRAPES-3 experimentis 8% and a comparatively short onset time of 10 h andthereafter the recovery phase of the Fd commenced. The
onset and recovery phases of this Fd display a rich time
structure. The tracking feature of the muon telescope has
shown a systematic time shift in the onset of the Fd on a

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PHYSICAL REVIEW D 74, 052003 (2006)

time scale of an hour. By exploiting the tracking feature of
the muon telescope for analyzing the rigidity of the incoming particles, the rigidity dependence of the magnitude of
the Fd has been estimated.
II. GRAPES-3 EXPERIMENTAL SYSTEMThe GRAPES-3 experiment is located at Ooty (11.4 N
latitude, 76.7 E longitude, and 2200 m altitude), a popularmountain resort town in southern India. The GRAPES-3 air
shower experiment has been designed to have one of themost compact configurations of the conventional type arrays with a separation of only 8 m between the adjacentdetectors which are deployed in a symmetric hexagonal
geometry. A schematic layout of the GRAPES-3 array isshown in Fig. 1. The observations were started in early
2000 with 217 detectors, located within the inner 8 rings
which are shown as filled circles in Fig. 1. Depending on
the supply of the plastic scintillators which are being
continuously fabricated in-house, additional detectors arebeing regularly added. Currently the experiment is being
operated with over 320 detectors. The final goal is to install
15 rings with 721 detectors, covering a total distance of
120 m from the center of the array. The detectors yet to beinstalled are shown as open circles in Fig. 1 [23].
A very large area tracking muon telescope operating aspart of the GRAPES-3 experiment [23,24] is an unique
instrument used to search for the high energy protons
emitted during the active phase of a solar flare. The
muon telescope is capable of providing a high statistics,

FIG. 1 (color online). A schematic layout for the 721 detector(open circles) GRAPES-3 shower array of which 217 detectors(filled circles) used at present are shown. Each of the 16 squares
represents a 35 m2 area muon tracking detector with E 
1 GeV.

directional study of muons. The GRAPES-3 muon telescope has a total area of 560 m2 , consisting of 16 modules,
each of which is 35 m2 in area. These modules are locatedvery close to each other as shown in Fig. 1. A cluster offour 35 m2 area neighboring modules, located inside acommon hall, constitutes one supermodule with a total
area of 140 m2 . In Fig. 2, a picture of the four halls, eachhousing one supermodule of the tracking muon telescope,is shown. Some of the scintillator detectors of the
GRAPES-3 air shower array (sealed inside the pyramid
shaped housing) which have been placed on the roof andaround these halls are also visible [23]. The energy threshold of the telescope is 1 GeV for the muons arriving along
the vertical direction. The cutoff rigidity at Ooty is 17 GV
in the vertical direction and varies from 12 to 42 GV acrossthe field of view of the tracking muon telescope as shown
later in Fig. 9.
The basic detector element of the muon telescope is arugged proportional counter (PRC) made from a 600 cm
long steel pipe with 2.3 mm wall thickness and a squarecross-sectional area of 10  10 cm2 . A muon detector
module with a sensitive area of 35 m2 consists of a totalof 232 PRCs arranged in four layers of 58 PRCs each, withalternate layers placed in orthogonal directions. Two successive layers of the PRCs are separated by a 15 cm thickreinforced cement concrete layer, consisting of 60 cm 
60 cm  15 cm blocks. The four-layer PRC configurationof the muon modules permits us to have a two-dimensional
reconstruction of the muon tracks in two vertical, orthogonal planes, namely, XZ and YZ planes, where the Z axis
lies along the vertical direction. The vertical separation of
the two layers of the PRCs in the same projection plane is50 cm, which allows the direction of the muon track to

FIG. 2 (color online). A view of the GRAPES-3 experimentshowing the four halls housing the four supermodules of themuon detector. The scintillator detectors of the GRAPES-3 air
shower array, placed on the roof and around the muon halls arealso visible.

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PHYSICAL REVIEW D 74, 052003 (2006)

be measured to an accuracy of 6 in each of the two
projected planes.
In order to achieve an energy threshold of 1 GeV for thevertical muons, a total thickness 550 g cm2 in the formof concrete blocks has been employed as the absorber. Thishas been done by placing a total of 15 layers of concrete
blocks above layer 1, as is shown schematically in Fig. 3. A
unique feature of the GRAPES-3 muon detector is therobust structure of the PRCs which has permitted it to
support the huge load of a 2.4 m thick concrete absorberin a self-supporting manner. The concrete blocks have beenarranged in the shape of an inverted pyramid to shield thePRCs, with a coverage of up to 45 for the incident muons.
The detectors have an energy threshold 1 ( sec) GeV forthe muons incident at a zenith angle of . The cross sectionof a muon detector module is shown schematically in
Fig. 3. As mentioned earlier, a cluster of four such modules
separated by a horizontal distance of 130 cm at the baseconstitutes one supermodule. The GRAPES-3 muon detector contains a total of four supermodules [24].
Because of the sensitivity of the PRCs to low-energy rays from the radioactivity present in the concrete absorber,
individual PRCs display sizable counting rates of200 Hz. When an output is generated, if any one of the
58 signals is present, is termed an OR. A coincidence of theOR outputs from the four layers of a module is used togenerate a 4-layer trigger. Despite the high counting ratesof the individual PRCs due to the radioactivity that ispresent in the surrounding absorber, the 4-layer coincidence triggers are relatively free from this backgroundand represent the passage of through going muons. The
observed 4-layer muon counting rate of 3200 Hz per
module yields a total counting rate 3  106 min1 forall 16 modules. This large rate of muons permits even asmall change of & 0:1% in the muon flux to be detected
over a time scale of 5 min after the application of appropriate corrections for the variation in the atmosphericpressure.
Unfortunately, during the period 28–29 October 2003,
only three out of four supermodules were continuously
operating. Therefore, in all subsequent discussions theLayer−3
Layer−2
Layer−1
Layer−0data from only 12 working modules have been used. In
Fig. 4, the percent variation in the 4-layer muon rate asrecorded by one of the 12 working modules relative to theremaining 11 modules is shown for 28 October 2003. The
horizontal axis in Fig. 4 represents the coordinated universal time (UTC) in hours for 28 October 2003. The percent
variation observed in the muon rates is consistent with astatistical error of 0.04%, expected from the total muoncounts (  6  106 ) in the sample. Similar results havebeen obtained for the remaining modules. However, alimiting value of the precision of 0:016% is obtained,
in cases where the statistical error from the total muoncount in the sample is significantly smaller than this value.
We attribute this limiting precision (0.016%) to the overall
stability of the detector system and to the irreducible
systematic errors present in the data. This high level ofprecision permits a very detailed study of the cosmic rayintensity variation to be carried out. For example, the onsetof the Forbush decrease on 29 October 2003 which started
20 h after the start of an X17 class solar flare at 09:51
UTC on 28 October 2003 [25] has been studied in considerable detail with the GRAPES-3 instrument.
However, most of the detected muons are generated by* 20 GeV galactic cosmic rays, which are of little interest
to us here, except for the fact that they form a stable anddominant background to the variation in the muon flux
produced by the solar flare. The statistical variations in
the muon rate are relatively small in magnitude because of
the large area of the GRAPES-3 muon telescope, where3  106 muons are detected every minute from the hits
recorded in the 3712 PRCs of the GRAPES-3 muon telescope [24]. The muon data is grouped online every 10 sec,
into solid angle bins of 0:05 sr, consistent with theangular resolution of the muon telescope as detailed below.
However, despite the drastic reduction in the data generated due to this angular segmentation, we still record
9 GB of data from the four supermodules every day.

Deviation from Mean



0.1 %

0%

-0.1 %-0.2 %
0

4

8

12

16

20

24

Time in hours (UTC)

FIG. 3 (color online). A schematic display of the 4-layer
tracking muon detector module with 58 PRCs per layer. The
four layers of the PRCs labeled Layer-0, Layer-1, etc. areembedded in concrete blocks. Inclined lines represent a set ofparallel muon tracks.

FIG. 4 (color online). The percent deviation in the 4-layer
muon counting rate of one of the 35 m2 area module relative
to the mean rate of remaining 11 modules as a function of time,on 28 October 2003. The rates are taken for 30 min intervals.

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PHYSICAL REVIEW D 74, 052003 (2006)

The relativistic solar particles can provide a key to theunderstanding of the acceleration processes at work in theSun. The global network of NMs has been used to detect
relativistic neutrons produced in the atmosphere from theinteractions of the high energy particles from the Sun,
which were accelerated during the solar flare on
28 October 2003 around 11:10 UTC nearly 80 min after
the detection of this flare in x rays. This solar flare had
followed in the wake of the highly disturbed interplanetaryconditions prevailing because of the earlier eruptions that
had occurred in the Sun.
We have searched for the direct emission of protons at
the time of this solar flare. However, during the period ofinterest as mentioned earlier, only three out of four supermodules were operating and therefore the total area of themuon detector was reduced to 420 m2 . In Fig. 5 thepressure-corrected GRAPES-3 data of 28 October 2003
for the muon rates summed over the three working supermodules is displayed. The variation in the rate of muons,
with an amplitude of 1% due to the diurnal anisotropy isclearly visible. However, no statistically significant excess
has been recorded in the time interval of 30 min after thestart of the solar flare at 09:51 h UTC, indicated by the firstarrow in Fig. 5 or after 11:10 h UTC, indicated by thesecond arrow when the global network of NMs started
recording increase in the rate of neutrons.
The value of 30 min has been selected as the energetic(multi-GeV) protons are expected to arrive within 1–2
times the light travel time (8–16 min) from the Sun to
the Earth as given in the model quoted above [3]. However,in the absence of the precise knowledge of the energy
spectrum of solar protons at high energy and that of theinterplanetary magnetic field, there is a considerable degree of uncertainty in the actual travel time of the energetic
solar protons. The NM data showed the presence of the
proton flux over an extended period of time. Thus an
39200

interval of 30 min is expected to provide a reasonable
time window for the detection of any direct high energy
protons from the solar flare. The exact choice of the time
window is not a critical parameter because the upper limiton the proton flux as is described below, is not particularly
sensitive to the duration of the time window selected.
To further investigate the possibility of emission of
protons on time scales shorter than 30 min, the data has
been analyzed for intervals of 1, 3, and 6 min for a limited
time window from 09:00 to 13:00 h UTC and is shown in
Fig. 6. The percent change in the total rate of muons,detected by the three working supermodules for the 1, 3,
and 6 min intervals is shown in the top, middle, and bottom
panels of Fig. 6, respectively. The arrow on the left in eachpanel indicates the start of the x-ray flare and the arrow on
the right the start of a flux detected by NMs. However, no
statistically significant excess is seen anywhere during theinterval 9–13 h UTC. The systematic decrease in the muonrate, seen in Fig. 6, is due to the well-known diurnal
anisotropy.
Since >1 GeV muons are secondaries produced by the
primary protons of energy * 20 GeV in the atmosphere,
these observations have been used to place a 99% C.L.
upper limit of 1:4  106 cm2 s1 sr1 on the wide-angle(  2:5 sr) flux of solar protons. The tracking muon tele-

0.6
0.4
0.2
0
3 Minute

0.8
0.6
0.4
0.2
0

39000
Muon Rate (s -1 )1 Minute

0.8

Rate Change (%)

III. SOLAR PROTON EMISSION

6 Minute

0.8
38800

0.6
0.4

38600

0.2

38400

0
38200
0

4

8

12

16

20

24

9

10

11
12
Time in Hours (UTC)13Time in hours (UTC)

FIG. 5 (color online). The variation in the total muon rate forthe three supermodules, for 28 October 2003 after atmosphericpressure correction. Start of the solar flare is indicated by the firstarrow. The second arrow indicates the start of the prompt
component detected by NMs.

FIG. 6 (color online). The percent variation in the total muon
rate for the three supermodules, for 28 October 2003 9–13 hUTC; (i) for 1 min interval (top panel), (ii) for 3 min (middle),and (iii) for 6 min (bottom). The left and right arrows in eachpanel indicate the start of the x-ray flare and that of the prompt
component detected by NMs.

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PHYSICAL REVIEW D 74, 052003 (2006)

scope is being operated round the clock to search for direct
protons during the solar flares which may occur in the
future.
IV. NARROW-ANGLE SOLAR PROTON EMISSION
It is well known that the solar protons of energies of tensof GeV, need not be observed from the direction of the Sun,due to their deflection in the interplanetary and geomagnetic fields. Therefore, it is possible to enhance the sensitivity of the muon detector by incorporating themeasurement of the arrival direction of the muons. This
novel feature has been incorporated into all four supermodules of the GRAPES-3 detector by recording the muon
directions into 225 solid angle bins. This capability has
been achieved after the installation of new, additional
direction-sensitive trigger electronics and a separate dataacquisition system for each of the four supermodules. As isshown in Fig. 7, for each PRC in the lower layer, the muon
angle is determined and binned into 15 angular bins basedon the specific PRC triggered in the upper layer fromamong the 15 PRCs, one directly above (central PRC)
and 7 each on either side of the central PRC. This angularbinning is carried out in each of the two orthogonal projections (XZ and YZ), thereby generating a twodimensional 15  15  225 solid angle map of the muon
directions. The contents of the 225 solid angle bins arerecorded, once every 10 s, thus permitting a continuous
monitoring of the directional flux of the muons in the sky.
Although the excess in the muon rate can be searched forin any of the 225 solid angle bins, it is expected that the fluxof the solar protons from a flare would be spread over
several bins. This directional spread could arise from anumber of factors, namely, (a) the size of the bin is comparable to the angular resolution of the muon telescope,(b) the magnetic field in the solar flare and in the interplanetary regions, (c) the terrestrial magnetic field, and(d) the transverse momentum imparted in the atmosphericcollisions. Therefore, for the purposes of the present study,
the muons have been regrouped into 3  3  9 coarse
solid angle bins, as shown schematically in Fig. 8. This
regrouping of the data has been done by either combining a
set of 3  5 or 5  5 solid angle bins. The exception being
the vertical direction which combines 3  3 bins. This
peculiar choice of binning is dictated by the fact that themuon flux is comparatively larger for the near central
directions (N, E, W, S) as compared to the outer directions(NE, SE, NW, SW). This regrouping of the data results in a
relatively similar solid angle coverage for the nine bins.
Therefore, this rebinning also results in muon statistics forvarious bins which are not too dissimilar from each other.
The geomagnetic cutoff rigidity for the field of view(FOV) of the GRAPES-3 muon telescope varies significantly for the nine solid angle bins. In Fig. 9 the contours ofconstant geomagnetic cutoff rigidity in the FOV are superimposed over a schematic of the nine solid angle bins ofmuon arrival directions in two orthogonal projected planes.
The geomagnetic cutoff rigidity varies from a low value of
12 GV in the West to 42 GV in the East direction within the
FOV of the GRAPES-3 tracking muon telescope. The
measurement of the direction of muons in GRAPES-3
experiment permits the effect of the rigidity on various
solar phenomena to be systematically studied as describedlater in Sec. VI.

7
6
5
4

−1 0

1

2

4

7

NW

N

NE

W

V

E

SW

S

SE

3
2
1
0
−1
−2
−3
−4
−5
−6
−7

54O

38O

21O 11O 0 O −11O

−7

FIG. 7. A schematic view of muon arrival angle selection
based on the PRC triggered in the lower and 15 PRCs directly
above, in the upper layer. Triggered PRCs are shown as filledsquares.

−6 −5 −4 −3 −2 −1

0

1

2

3

4

5

6

7

FIG. 8. A schematic of the nine groups of muon arrival directions which combine to form solid angle bins; 3  3 vertical bin
V or four 3  5 central bins N, E, W, S or four 5  5 outer bins
NE, SE, SW, NW.

052003-7

T. NONAKA et al.

PHYSICAL REVIEW D 74, 052003 (2006)

FIG. 9. A schematic of the nine solid angle bins of muon
arrival directions. Also shown are contours of constant geomagnetic cutoff rigidity in the field of view. The cutoff rigidity varies
from 12 to 42 GV in the FOV of GRAPES-3 muon telescope.

With the enhancement in the directional sensitivity,
achieved through this angular segmentation, the fractionalvariation in the muon counting rate for time intervals of
6 min is shown in Fig. 10 for the nine solid angle bins. The

diurnal anisotropy observed in the wide-angle data shownin Fig. 5 is also visible for all nine directions. However, thetime of the maximum amplitude of the anisotropy seems tochange with the arrival direction. For example, the maximum amplitude is reached earliest for the three East (NE,E, SE) directions, followed by the Center (N, V, S) andfinally the West (NW, W, SW) directions. This time evolution is discussed, in greater detail later on, in Sec. V.
No statistically significant excess is seen, from any of
the nine directions in the muon rate, during an interval of
30 min, after the start of the solar flare at 09:51 h UTC on
28 October 2003. A 99% C.L. upper limit of 4 
106 cm2 s1 sr1 on the narrow-angle (  0:3 sr) fluxof the solar protons of * 20 GeV from each of the nine
angular regions has been placed. An analysis of this data on
shorter time scales, similar to that of the wide-angle data,
described earlier in Sec. III, also did not reveal any statistically significant excess on time scales of 1, 3, and 6 minfor all nine directions.
The GRAPES-3 wide-angle upper limit along with thoseobtained for the solar flares of 29 September 1989 by the
Baksan group [15] and of 14 July 2000 by the L3 group[17] are also shown in Fig. 11. A flux of neutron of 2 104 cm2 s1 above 10 GeV claimed by the AGASA
group [16] for the solar flare of 4 June 1991 is shown in
Fig. 11. For a comparison, the flux of particles recorded
using NMs during one of the largest flares recorded on
23 February 1956 and the one on 15 April 2001 are also
shown in Fig. 11.

2

10 2
GOES

Integral Flux (cm-2 s -1sr -1)

NE

0

E

-2
Rate Change (%)

SE
-4

N

-6

V
S

-8

1
NM

10

-2

10-4

NW

-10

W

10-6

SW

-12

0.1

1

10

100

1000

Energy (GeV)

-14

0

4

8
12
16
Time in Hours (UTC)20

24

FIG. 10 (color online). The observed variation in the muon
rate on 28 October 2003 for nine solid angle bins is shown, for
6 min intervals. Successive plots have been shifted down by
1.5% for clarity. The start of solar flare is indicated by the arrow.

FIG. 11 (color online). Integral flux of protons with energy on
28 October 2003; (a) based on 100 –600 MeV GOES-10/11 andsub-GeV balloon data, (b) worldwide NM data, (c) 䊏 * 20 GeV
GRAPES-3 upper limit, (d) 䉱 >10 GeV AGASA flux on 4 June
1991, (e) 䉲 >40 GeV L3 upper limit on 14 July 2000, (f) 夹>500 GeV Baksan upper limit on 29 September 1989, (g) 䊉NM data on 15 April 2001, (h) 䉬 NM data on 23 February 1956.

052003-8

DID THE 28 OCTOBER 2003 SOLAR FLARE . . .

PHYSICAL REVIEW D 74, 052003 (2006)

The GOES-10/11 instruments had recorded a rapid increase in the flux of >100 MeV proton at 11:30 h on
28 October 2003. The flux of proton at >100 MeV had
increased by almost 2 orders of magnitude in the next 2 h
and then started declining slowly 12 h later around 23:30 h
UTC on 28 October 2003. The straight line at the top left ofFig. 11 represents a power law fit to the integral energy
spectrum of the protons, in the energy range 100 –600 MeV as measured by the GOES-10/11 satellites on
28 October 2003 [26]. The power law fit to the energy
spectrum, extrapolated beyond observed GOES energies,
is shown by a dotted line. The measurements made on theflux of sub-GeV protons on a balloon flight on the same dayis also in excellent agreement with the GOES data [3]. The
second straight line shown in Fig. 11 represents a powerlaw fit to the integral energy spectrum of GeV protons
based on the measurements of the worldwide network of
NMs on 28 October 2003. The power law fit to the energy
spectrum, extrapolated beyond observed NM energies, isshown by a dotted line.
In the absence of the NM data, the upper limit from the
GRAPES-3 experiment would have been consistent withthe extrapolation of the low-energy GOES data, as seen in
Fig. 11. However, the NM data of 28 October 2003 displays an energy spectrum, which seems to significantly
harden in the GeV energy range. However, this trend
must reverse itself, and the energy spectrum beyond
5 GeV should either soften rapidly or there should be acutoff in order to be consistent with the upper limit at
20 GeV derived from the GRAPES-3 data. However, in
the absence of rapid softening of the energy spectrum or acutoff, the extrapolation of the NM data is 2 orders ofmagnitude higher than the GRAPES-3 upper limit. The
lower limit on an assumed power law spectral slope basedon the NM and GRAPES-3 data is   5:3. If the flux of
20 GeV protons were comparable to the positive flux
values reported by the AGASA and other groups, the
GRAPES-3 would have detected it as an increase * 5%of the ambient cosmic ray flux. Such a large increase would
be almost impossible to miss in the GRAPES-3 data.
Here note that the muons detected by the Baksan and theL3 groups were at energies which were significantly higher(1–2 orders of magnitude) as compared to the GRAPES-3
experiment. However, the energy threshold of solar cosmic
rays responsible for the muons detected by the AGASA
group was much closer to that of the GRAPES-3 experiment and, therefore, permitted a more meaningful comparison of their results. However, unlike the GRAPES-3
experiment, the absence of a tracking muon telescope in
the AGASA experiment compelled that group to utilize
their unshielded plastic scintillator detectors in the array,for the measurement of the muons produced in the atmosphere by the solar flare particles. Therefore, the AGASA
group had to rely on the net increase in the counting rate of
their scintillation detectors, which then was attributed tothe muons produced by the solar flare particles. A clear cut
identification of the muons and their arrival direction in a
tracking muon telescope as in the case of the GRAPES-3
was not possible for the AGASA experiment.
Since the 4 June 1991 solar flare occurred near the east
limb of the Sun, the protons originating there are expected
to arrive at the Earth after a delay of 1–2 h after the onsetof the x-ray emission based on earlier flare observations[16]. However, the near coincident (within a few minutes)
increase in the muon flux with the onset of the x-ray
emission prompted the AGASA group to conclude thatthe parent particle of muons should be neutrons. The
presence of the three muon detectors at altitudes of2700 m (Mt. Norikura), 900 m (AGASA), and sea level(Nagoya) had permitted them to estimate the attenuation
length of the neutrons in the atmosphere. The measured
attenuation length turned out to be anomalously large at
200 g cm2 as compared to the normal value100–120 g cm2 . The AGASA group also estimatedthat the spectral index of the power law describing theneutron flux was 4–5. The flux of the solar neutrons forthe 4 June 1991 flare was estimated to be 1–2 104 cm2 s1 which is about 2 orders of magnitude larger
than the 99% C.L. upper limit of 1:4  106 cm2 s1 sr1on the wide-angle flux of the particles placed by GRAPES3 for the 28 October 2003 flare. In addition, the tracking
feature of GRAPES-3 muon telescope allowed us to placea slightly less stringent 99% C.L. upper limit of 4 
106 cm2 s1 sr1 on the narrow-angle flux of solar flareparticles, for which, in absence of tracking, there is nocounterpart in the AGASA experiment.
Based on the past observations, the energy spectra of theparticles accelerated in the solar flares display considerable
complexity. It is reasonable to assume that the energy
spectrum of the particles accelerated during a solar flare
follow a power law. A compilation of the major solar flare
events during the past 60 yr has shown that the spectral
index of the assumed power law shows progressive steepening with increasing energy of protons, reaching a valueof 3.5– 4.0 at energies above 10 GeV [27]. However, it isvery difficult to arrive at a definitive conclusion on theenergy spectrum of the solar flare particles in view of thehuge variation seen from flare to flare; based on the pastclaims, it is likely that the solar flares may accelerate
protons to energies well above 100 GeV [12,15]. A lower
limit on the spectral index of an assumed power law for theenergy spectrum of the 28 October 2003 flare can be placedat   5:3 based only on the worldwide NM and the
GRAPES-3 data. However, in view of the rapidly changing
spectral indices and/or cutoffs observed in the case of mostof the solar flares, the actual value of the spectral index islikely to be significantly higher [27]. The narrow-angle
upper limits on the flux of solar flare particles derived forthe nine bins are very similar and for that reason a common
value of 4  106 cm2 s1 sr1 has been quoted above.

052003-9

T. NONAKA et al.

PHYSICAL REVIEW D 74, 052003 (2006)

V. FORBUSH DECREASE ON 29 OCTOBER 2003

Rate Change (%)

0
1σ
-2

-4

-6

-8
0

4

8

12

16

20

24

Time in hours (UTC)

FIG. 12 (color online). The observed variation in the muon
rate on 29 October 2003 from the vertical solid angle bin ‘‘V.’’
Muon rates shown are for 6 min intervals. The diameter of thesmall dot to the right of legend 1 equals 1 sigma statistical error
in data.

the neutron data, the hourly rate of change of the neutron
flux for the Tibet NM [30] is shown in Fig. 13 as a functionof time in hours UTC on 29 October 2003. The geomagnetic cutoff rigidity of 14 GV in the vertical direction at
Tibet is only slightly lower than 17 GV at Ooty. The
proximity of the longitudes of Ooty and Yangbajing in
Tibet (only 13.8 apart) implies that the same region of
the sky is observed at these two locations almost concurrently. The onset of the Fd is clearly visible in Fig. 13
around the same time as in GRAPES-3 muon data shown in
Fig. 12. The change in the neutron count rate has been
calculated after normalizing it to the value observed at
00:00 h UTC on 29 October 2003. The magnitude of the

Tibet NM
0

Rate Change (%)

In addition to the acceleration of particles due to theenergy released from the coronal loops, the solar flares arealso associated with the coronal mass ejections which areblobs of plasma ejected during the flaring activity. An
expanding CME-driven shock front accelerates ions at
various helio-longitudes. Apart from the complex structureof a CME and its interaction with the heliospheric magnetic cavity, its effect on the structure of the geomagnetic
field is also very important for various studies at the Earth.
It is widely recognized that the CMEs and high-speed
streams from the coronal holes have considerable influenceon the propagation of the particles within the heliosphere.
The interaction of these flows with the quiet solar windcreates regions of compressed heated plasma and shocks
which are responsible for further acceleration and modulation of the particles incident over the atmosphere of the
Earth. Detailed observations on the modulation of thecosmic rays, using NMs and muon telescopes, provide
important information for modeling the propagation andacceleration of the solar energetic particles in the interplanetary medium. The short-term modulations such as Fds
may be related, either directly to the material ejected fromthe Sun or to the transient interplanetary shocks generated
by these ejections [19,20,28].
The capability to observe the directional muon flux, withthe three working supermodules, provides a completely
new view of the evolution of a Fd, which should help inmodeling the propagation of the high energy particles in
the interplanetary and in the near-Earth regions. This may
be seen from the detailed information available on the Fd,
observed on 29 October 2003 from the nine solid anglebins, each with a solid angle coverage of 0:3 sr as displayed in Fig. 8. The observed muon rate for time intervalsof 6 min for one of the nine bins, namely, ‘‘V,’’ after
applying appropriate corrections for the variation in theatmospheric pressure as a function of time in hours UTC on
29 October 2003, is shown in Fig. 12. The rapid onset of the
Forbush decrease starting around 7 h and lasting until 18 h
UTC is clearly visible. Also shown in Fig. 12 is a dot next
to the legend 1 whose diameter represents the size of a 1error bar. Therefore, the rich structure in the muon rate
seen during the onset of Fd is to be regarded as real, and it
indicates the rapid variation in the evolution of the Fd on atime scale of minutes. It has been pointed out that a ‘‘twostep’’ structure observed in several Fds is caused by theejecta and the shock effects of the CME [28]. Therefore,
any model proposed to account for the particle accelerationand propagation through the interplanetary space during asolar flare should also be able to explain this rich time
structure during the onset phase of the Fd.
The Tibet neutron monitor (NM) operated by the JapanChina collaboration is located at Yangbajing in Tibet(30.1 N latitude, 90.5 E longitude and 4300 m altitude)[29]. For a comparison of the GRAPES-3 muon data with

2

-5

-10

-15
0

4

8
12
16
Time in hours (UTC)20

24

FIG. 13 (color online). The hourly variation in the observed
neutron rate on 29 October 2003 for the Tibet neutron monitor.
The change in rate is normalized, relative to the rate at 0 h UTC.

052003-10

DID THE 28 OCTOBER 2003 SOLAR FLARE . . .

PHYSICAL REVIEW D 74, 052003 (2006)

Fd is 15% as compared to 8% seen in GRAPES-3
muon data. The profile of the onset of the Fd is measured
with much greater precision in the GRAPES-3 muon data,
due to a higher count rate and shorter sampling interval.
Therefore, the rich time structure seen in Fig. 12 is far lessconspicuous in the NM data shown in Fig. 13. However, theoverall profile of the Fd appears to be rather similar in bothcases with a double valley structure near the bottom.
The observed muon rates for time intervals of 6 min for
the nine bins, after applying the appropriate corrections forthe variation of the atmospheric pressure, are shown in
Fig. 14. Note that there is a systematic variation at thetime, when a specific feature in the muon rate appears in agiven solid angle bin. To highlight this aspect, the data in
Fig. 14 are displayed after grouping them in three separate
sets labeled ‘‘East,’’ ‘‘Center,’’ and ‘‘West,’’ respectively.
The east set contains the data from the three eastern solid
angle bins, namely, NE, E, and SE. Similarly the center set
contains the data from solid angle bins N, V, S, and finally
the west set contains data from the solid angle bins NW, W,
SW. For clarity, data shown in Fig. 14 for center and west
sets have been shifted downwards by 4% and 8% relativewith respect to the east set. The three arrows in Fig. 14
point to the time when the Fd reaches its minimum value
for each of the three sets.
To get a more quantitative comparison of the profiles of
the Fd as a function of muon arrival direction, we havecarried out a systematic correlation study between the data

Ci;j s  Ri t Rj t

100

-5

-10

-15

0

4

8
12
16
Time in Hours (UTC)20

24

FIG. 14 (color online). The observed variation in the muon
rate on 29 October 2003 from the nine angular bins grouped into
three sets, namely, east (NE, E, SE), center (N, V, S), and west
(NW, W, SW). Muon rates shown are for 6 min intervals.
Successive sets have been shifted downwards by 4% for clarity.
Each pointer represents the time of the Fd minimum for the
3 sets, respectively.

s =Ri t 2 Rj ts 2 0:5 ; (1)where Ri is the percent change in muon rate observed forsolid angle bin i, and t runs over the period 7–19 h UTC on
29 October 2003 which coincides with the onset of the Fd,and the time-shift ‘‘s’’ is varied from 4 to 4 h in stepsof 0.1 h. The paired data for s  0 shows a very highdegree of correlation of 99:7  0:2 %, 98:8  0:2 %,and 96:3  0:3 % for the categories 1, 2, and 3,
respectively. As expected the correlation coefficient gradually decreases with the increasing angular separation in theeast-west direction. However, there is no such prominent
effect in the north-south direction. In Fig. 15 a plot of thecorrelation coefficient as a function of time-shift s is shownfor one pair, from each of the three categories. The correlation coefficient shows a maximum at a time shift of zero,
0.6 and 1.3 h, respectively, for the three categories.
The sensitivity of the GRAPES-3 muon data is highlighted by the fact that the correlation coefficient for
category-2 data for the pair ‘‘N’’ and ‘‘NE’’ systematicallyincreases from 98.8% to 99.6% as the time shift was varied
from 0 to 0.6 h in the correct direction as seen from Fig. 15.
Although the increase of 0.8% in the magnitude of thecorrelation coefficient is modest, given the small value of

Correlation Coefficient (%)

Rate Change (%)

0from a pair of solid angle bins. For this purpose the paired
data is classified into three categories, namely, (1) where
the two solid angle bins in the pair belong to the same set (atotal of 9 pairs), (2) where the solid angle bins belong to theadjacent sets (a total of 18 pairs), and (3) where the solid
angle bins belong to the extreme sets, namely, one eachfrom east and west sets (a total of 9 pairs). The correlation
coefficient Ci;j s for the paired data of two solid angle binsi and j is calculated as a function of the shift in time ‘‘s’’ asshown below,

a

b

c

99
98
97
96
95
-4

-3

-2

-1
0
1
Time Shift (hours)

2

3

4

FIG. 15 (color online). Observed variation in the correlation
coefficient with time shift (hours) for pairs of angular bins; (a) category-1, W-SW, (b) category-2, N-NE, (c) category-3, E-W
for the data of 29 October 2003.

052003-11

T. NONAKA et al.

PHYSICAL REVIEW D 74, 052003 (2006)

statistical error (  0:1%) in the GRAPES-3 data, even thismodest increase is significant. But a time shift of 0.6 h in
the opposite direction decreases the correlation coefficient
to a comparatively lower value of 97%. The observed
increase in the magnitude of the correlation coefficient of
3.1% is much larger for the category-3 data, for the pair‘‘E’’ and ‘‘W’’ from 96.3% to 99.4% as the time shift isincreased from 0 to 1.3 h in the correct direction as seenfrom Fig. 15. However, a time shift of 1.3 h in the opposite
direction decreases the correlation coefficient to a value
well below 90%.
The angular resolution of the muon telescope is 6 ineach of the two vertical projection planes and the solid
angle covered by each of the nine bins is 0:3 sr.
Therefore, the effective hour angle spanned by each bin
is 25 . Thus it should take about 1.7 h for a source in thesky to transit across the field of view of each of the ninesolid angle bins. Therefore, purely from the hour angle
arguments, it is expected that a Fd or any other phenomenon induced by the Sun, which produces a stable anisotropy should exhibit a time lag 1:7 h for the category-2
and 3:3 h for the category-3 data.
To test this ‘‘hour angle’’ hypothesis we have carried out
an identical investigation of the correlation in the data of
27 October 2003. This day was selected as it was 1 dayprior to the occurrence of the solar flare. The correlation
coefficient is calculated for 36 pairs of solid angle bins on
27 October 2003, over a duration of 12 h which includes
the diurnal anisotropy, using Eq. (1). This was done inexactly the same manner as for the data of 29 October
2003. In Fig. 16 a plot of the correlation coefficient as afunction of the time-shift s using Eq. (1) is shown for thethree pairs, one each from the three categories. The threepairs selected are W-SW, N-NE, and E-W which are iden100
Correlation Coefficient (%)

a

b

c

tical to the pairs used for the data of 29 October 2003 asshown in Fig. 15. The data for s  0 still shows a sizable
degree of correlation; (i) 96:4  1:7 % for the 9 pairs of
the category-1, (ii) 91:0  2:9 % for 18 pairs of thecategory-2, and (iii) 74:5  5:9 % for the 9 pairs of
category-3, respectively.
The correlation coefficient marginally increases from
96.4% to 96.8% for a time shift of 0.32 h with a dispersionof 0.28 h for the category-1 data. The correlation coefficient increases from 91.0% to 97.0% for a time shift of
1.92 h with a dispersion of 0.31 h for the category-2 data
and from 74.5% to 95.3% for a time shift of 3.86 h with adispersion of 0.21 h for the category-3 data. Therefore, a
net time shift of 1:6  0:4 h is observed between the
category-1 and category-2 data which is very similar to
the expected value of 1.7 h based on the hour angle
calculations for the GRAPES-3 FOV. Similarly, the net
time shift of 3:5  0:4 h is observed between the
category-1 and category-3 data, which is also very similar
to the expected value of 3.3 h based on hour angle calculation for the GRAPES-3 FOV as seen in Fig. 16.
To summarize, the correlation study has shown a timeoffset of 0.6 and 1.3 h for categories-2 and -3 data, respectively, relative to the category-1 data during Fd of
29 October 2003. This is in stark contrast to a time offset
of 1.6 and 3.5 h for the categories-2 and -3, respectively,
relative to the category-1 data on 27 October 2003. As was
already explained, the 27 October 2003 values are inagreement with the expectations based on the hour angleof the FOV of GRAPES-3 tracking muon telescope.
Clearly, a rapid time evolution in the onset of the Fd on atime scale of 1 h across a field of view of 25 isrequired to explain this reduction in the time offset from
1.7 to 0.6 h and from 3.3 to 1.3 h for the category-2 andcategory-3 data, respectively. The physical interpretationof the observed time evolution would require a detailed
phenomenological modeling of the passage of the shock
front produced by the CME and the effects on the cosmic
ray intensity around the Earth.

95

VI. RIGIDITY DEPENDENCE OF THE
MAGNITUDE OF THE FORBUSH DECREASE OF
29 OCTOBER 2003

90

85

80
-2

-1

0

1
2
3
Time Shift (hours)

4

5

6

FIG. 16 (color online). Observed variation in the correlation
coefficient with time shift (hours) for pairs of angular bins; (a) category-1, W-SW, (b) category-2, N-NE, (c) category-3, E-W
for the data of 27 October 2003, 1 day before the solar flare.

As shown in Fig. 9, the geomagnetic cutoff rigidity forthe nine solid angle bins in the field of view of the
GRAPES-3 tracking muon telescope varies from 12 GV
to 42 GV. This fact combined with the high quality
GRAPES-3 muon data, makes it possible to study therigidity dependence of the magnitude of the Forbush decrease observed on 29 October 2003. The mean geomagnetic cutoff rigidity for the nine solid angle bins varies
from 14.3 GV in west to 24.0 GV in the east direction. The
muon rates averaged over 6 min successive intervals for thenine solid angle bins are used to estimate the magnitude of
the Fd in each direction. In Fig. 17(a) the magnitude of the052003-12

DID THE 28 OCTOBER 2003 SOLAR FLARE . . .

PHYSICAL REVIEW D 74, 052003 (2006)

Fd is shown, as a function of the geomagnetic cutoff
rigidity for the nine bins. A clear inverse correlation between the geomagnetic cutoff rigidity and the magnitude of
Fd can be noticed, because the magnitude of the Fd decreases as the cutoff rigidity increases. A change of 1.6% in
the magnitude of the Fd is seen in going from 14.3 to24 GV. However, another interesting feature of this variation is the setwise clustering of the values of the Fdmagnitudes into groups of three bins (the three sets aredescribed in Sec. V). The three east bins labeled NE, E, SE
and represented by filled circles in Fig. 17(a) form acluster. But the middle bin ‘‘E’’ shows a sizable departurein its magnitude (  0:4%) as compared to the two outerbins NE and SE. Similarly the three ‘‘center’’ bins labeled

Magnitude of Forbush Decrease (%)

9
V

W

(a)
8.5
S
N
NW

8
SW

E

7.5
SE
NE

Ar  K%  r=30 

7Magnitude of Forbush Decrease (%)

N, V, and S and represented by filled triangles also show
clustering with bin V separated by 0:5% from the othertwo bins N and S. The displacement of bin W in the west
set, represented by filled squares, is the largest 0:8% withrespect to the remaining two bins, namely, NW and SW.
The inverse correlation between the magnitude of the Fdand the geomagnetic cutoff rigidity is an established feature of the phenomenon of Fd, and it can be readily understood on the basis of the rigidity spectrum of themodulation of the cosmic rays. The rigidity dependenceof the magnitude of the Fds have been measured by anumber of groups primarily using NMs. The rigidity dependence can be reasonably approximated by a power lawof the type K  r where r is the rigidity in GV and themeasured spectral index  ranges from 0.4 to 1.2 [28].
Although the observed displacement of the middle bin
relative to the remaining two bins in each set is rather small(  0:5%), given the small value of the statistical error in
our measurements this effect becomes significant. Notethat the middle bins (E, V, and W) in the respective setsof solid angle bins point to a direction, closest to the local
vertical direction. Therefore, not surprisingly, these threebins record a larger flux of muons due to a lower thresholdenergy, caused by a relatively smaller angle of incidence.
In order to determine the dependence of the magnitude‘‘Ar ’’ of the Fd (in %) in a given bin, Monte Carlosimulations have been carried out with an assumed rigidity
dependence described by the following power law:

W

(b)

V

8.5
SW

S
NW
N

8
E

7.5
SE
NE

7

14

16

18

20

22

24

Cutoff Rigidity (GV)

FIG. 17 (color online). (a) Observed variation in the magnitude of Forbush decrease (%); (b) simulations for the magnitude
of Forbush decrease (%). Both as a function of geomagnetic
cutoff rigidity (GV) for the three sets of nine solid angle bins,
namely, (i) NE, E, SE (filled circles), (ii) N, V, S (filled
triangles), (iii) NW, W, SW (filled squares). The middle bins
(E, V, W) display sizable deviations 0:5% from the remaining
two bins. Simulations are for a rigidity dependence of
Fd magnitude  12:8%  r=30 0:53 .

(2)

where r is the rigidity in GV. In simulations effects of the
instrumentation such as the rigidity cutoff, solid angle ofacceptance, angle of incidence, etc. are included to compute the expected magnitude of the Fd, for the nine solid
angle bins. A chi-square minimization method yielded thefollowing values for the parameters of the power law fit:K  12:8  0:3 %and  0:53  0:04.
This
Monte Carlo simulation gave a rms spread of 0.19% in
the magnitudes of observed and simulated data. A rms
value of 0.19% between observed and simulated data
sets, although larger than the intrinsic statistical error in
the observed data of 0:1%, the agreement between theobservations and the simulations is considered satisfactoryin view of the numerous simplifications made in carryingout the Monte Carlo simulations.
The simulated magnitudes of the Fd for the nine solid
angle bins are shown in Fig. 17(b). The simulated data
reproduces the pattern of Fd magnitudes which is very
similar to that of the observed data. The close agreement
between Monte Carlo simulations and observations indicates that the major features of the data have been accounted for in terms of known effects such as the angle
of incidence, energy threshold, etc. The value of the spectral index   0:53  0:04 of the rigidity dependence ofthe magnitude Ar of the Fd is in agreement with the othermeasurements (0.4 –1.2) as mentioned earlier [28]. Present

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PHYSICAL REVIEW D 74, 052003 (2006)

observations also demonstrate the ability of the GRAPES-3
muon telescope to make precision measurements on thesolar phenomenon such as the Forbush decrease, including
its rigidity dependence and anisotropy in real time.

over a relatively narrow range of geomagnetic cutoff rigidity (14.3–24.0 GV) has been derived using the datafrom the tracking muon telescope. The estimated valueof the spectral slope   0:53  0:04 is in agreementwith the earlier measurements which range from 0.4 to 1.2.

VII. SUMMARYACKNOWLEDGMENTS

In summary, the salient features of the tracking muon
telescope of the GRAPES-3 experiment at Ooty have been
described in relation to the search for the direct emission of
protons of * 20 GeV in association with the solar flare of
28 October 2003. A 99% C.L. upper limit on the wideangle flux of the solar protons is placed at 1:4 
106 cm2 s1 sr1 . This upper limit is significantly lower
than the limits and/or fluxes reported by the other experiments at comparable energies. The tracking feature built
into the GRAPES-3 muon telescope has been exploited tocarry out a narrow solid angle search of the direct protons.
The negative result obtained has been used to place a 99%
C.L. upper limit of 4  106 cm2 s1 sr1 . The significance of this novel feature of the GRAPES-3 muon telescope has been discussed in relation to the observations
pertaining to the propagation of the solar flare particles in
the interplanetary medium near the Earth. New information on the time evolution during the onset of the Forbush
decrease of 29 October 2003, on a time scale of an hour, isprovided. Rigidity dependence of the magnitude of the Fd,

We thank D. B. Arjunan, K. Manjunath, B. Rajesh, andC. Ravindran for their help in the testing, installation, andoperation of the proportional counters and the associatedelectronic modules. The administrative services of V.
Viswanathan during all phases of the experiment are gratefully acknowledged. It is a pleasure to thank A. A. Basha,
G. P. Francis, I. M. Haroon, V. Jeyakumar, and K.
Ramadass for their help in the fabrication, assembly, andthe installation of various mechanical components anddetectors. The Japanese members of the GRAPES-3
Collaboration acknowledge the partial financial support
from the Ministry of Education and Science of theGovernment of Japan for the GRAPES-3 experiment.
The GRAPES-3 Collaboration also thanks N. K. Mondal
and his colleagues in the TIFR-OCU Proton Decay
Collaboration for the loan of the proportional counterswhich have been used in the muon detector. Finally, we
thank the two anonymous referees whose critical comments led to a significantly improved presentation.

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