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Astronomy & Astrophysics manuscript no. 0666 February 2, 2004
(DOI: will be inserted by hand later)
Signature of oscillations in coronal bright points
I. UgarteнUrra 1 , J. G. Doyle 1 , M. S. Madjarska 1,2 , and E. O'Shea 3
1 Armagh Observatory, College Hill, Armagh BT61 9DG, N. Ireland
2 Mullard Space Science Laboratory, University College London, Holmbury St. Mary, Dorking, Surrey RH5 6NT, UK
3 Instituto de AstrofТРsica de Canarias, C/ vТРa LТactea s/n, 38200 La Laguna, Tenerife, The Canary Islands, Spain
Received / Accepted
Abstract. A detailed study of two consecutive bright points observed simultaneously with the Coronal Diagnostic Spectrometer
(CDS), the Extreme ultraviolet Imaging Telescope (EIT) and the Michelson Doppler Imager (MDI) onboard the Solar and
Heliospheric Observatory (SOHO) is presented. The analysis of the evolution of the photospheric magnetic features and their
coronal counterpart shows that there is a linear dependence between the EIT Fe ### 195 Љ flux and the total magnetic flux of
the photospheric bipolarity. The appearance of the coronal emission is associated with the emergence of new magnetic flux
and the disappearance of coronal emission is associated with the cancellation of one of the polarities. In one of the cases the
disappearance takes place #3н4 hours before the full cancellation of the weakest polarity.
The spectral data obtained with CDS show that one of the bright points experienced short time variations in the flux on a time
scale of 420н650 seconds, correlated in the transition region lines (O # 629.73 Љ and O ### 599.60 Љ) and also the He # 584.34 Љ
line. The coronal line (Mg ## 368.07 Љ) undergoes changes as well, but on a longer scale. The wavelet analysis of the temporal
series reveals that many of these events appear in a random fashion and sometimes after periods of quietness. However, we have
found two cases of an oscillatory behaviour. A subнsection of the O # temporal series of the second bright point shows a damped
oscillation of five cycles peaking in the wavelet spectrum at 546 seconds, but showing in the latter few cycles a lengthening of
that period. The period compares well with that detected in the S ## 933.40 Љ oscillations seen in another bright point observed
with the Solar Ultraviolet Measurements of Emitted Radiation (SUMER) spectrometer, which has a period of 491 seconds.
The derived electron density in the transition region was 3 Ѕ 10 10 cm -3 with some small variability, while the coronal electron
density was 5 Ѕ 10 8 cm -3 .
Key words. Sun: oscillations -- Sun: corona -- Sun: transition region -- Sun: UV radiation -- Sun: magnetic fields
1. Introduction
Coronal EUV/Xнray bright points (hereafter BPs) are small (20
н 30 Mm) coronal features of enhanced emission most easily
observed in the quiet Sun and coronal hole regions. They were
first observed and studied in soft Xнray rocket images (Vaiana
et al. 1970; Nolte et al. 1979) where they present an averн
age lifetime of #8 hours (Golub et al. 1974). However, they
have also been observed in radio (Marsh et al. 1980) and EUV
(Habbal & Withbroe 1981) where the emission has an average
lifetime of #20 hours (Zhang et al. 2001). The appearance and
hence their description is subject to the spatial resolution of the
instrument, so they are commonly presented as di#use clouds
with a central bright core of #10 Mm as seen in Skylab Xнrays
(Golub et al. 1974) or SOHO/EIT extremeнultraviolet images.
High spatial resolution images (1 ## н 2 ## ) have shown that BPs
consist of several miniature loops (Sheeley & Golub 1979) and
have a morphology similar to larger scale coronal structures
(Frank & Slater 2002).
Send o#print requests to: I. UgarteнUrra (iuu@star.arm.ac.uk),
http://star.am.ac.uk/preprints/
BPs are associated with photospheric bipolar magnetic feaн
tures (Krieger et al. 1971), with up to 2/3 of them being assoн
ciated with chance encounter and cancellation of preнexisting
magnetic features rather than the emergence of new magnetic
flux (Harvey 1985, 1993; Webb et al. 1993; Longcope et al.
2001). This process normally takes place at the network boundн
aries of superнgranular cells (Egamberdiev 1983; Habbal et al.
1990; Madjarska et al. 2003).
One of the main characteristics of BPs is their intensity
variability, as several studies in EUV spectral lines have shown.
Sheeley & Golub (1979) found that the constituent loops could
evolve on a time scale of #6 min. Habbal & Withbroe (1981)
and Habbal et al. (1990) using the Harvard experiment aboard
Skylab showed that they exhibit large variations in the emisн
sion of chromospheric, transition region and coronal lines, and
no regular periodicity or obvious correlation between the di#erн
ent temperatures was found. The time scales of these variations
were as short as the temporal resolution of the observations (5.5
minutes). Similar variations were found in Xнrays (Nolte et al.
1979). More recently, Madjarska et al. (2003) using SUMER
observations with a temporal resolution of 50 seconds, have

2 I. UgarteнUrra et al.: Signature of oscillations in coronal bright points
Fig. 1. BP2 location and semblance. Left image: context EIT Fe ### 195 Љ full disk image (8 November 2002, 00:36 UT) with an arrow
highlighting the location of the BP. Top right images: comparison of a closeнup of the feature in EIT and TRACE Fe ##/# 171 Љ taken 2 minutes
earlier. Bottom right: magnetic field contours (‘25 G, solid/dotted contours) overlaying the TRACE image. The arcade of loops resolved by
TRACE connects opposite polarities.
confirmed the presence within a BP of smallнscale transient
brightenings of the order of minutes. Nevertheless, concerning
BPs, very little has been done up to now to exploit the capabilн
ities of the CDS and SUMER spectrometers onboard SOHO.
Most of the studies have used EIT and MDI (also the Transition
Region And Coronal Explorer, TRACE) and have focused on
trying to understand the relationship between the BP coronal
emission and its magnetic counterpart (see for example PreТs &
Phillips 1999; Zhang et al. 2001; Longcope et al. 2001; Brown
et al. 2001). Our intention in this paper is to fill part of this
gap by studying the intensity variations of two BPs in di#erent
spectral lines with di#erent formation temperatures. We also
present the first wavelet analysis for BPs with the intention of
checking the apparent lack of periodicity found so far.
2. Observations
All along this paper we are going to refer to two consecutive
BPs observed in a coronal hole in November 2002. BP1 obн
served with EIT and MDI and BP2 observed as well with CDS.
2.1. CDS
On 7 November 2002, between 16:00 UT and 23:30 UT, and
pointing at BP2 at coordinates (--107 ## , +355 ## ), two studies reн
presenting temporal series of 81 and 30 seconds exposure time
were run with the Normal Incidence Spectrometer NIS/CDS
(Harrison et al. 1995) using the 4 ## slit and scan mirror tracking.
Telemetry restrictions only allow extraction of selected spectral
windows from the two NIS detectors. We selected two wide
windows (333.9 -- 372.6 Љ, 567.5 -- 631.3 Љ) for the first study,
in order to better remove the continuum contribution. In the
second study, we used four narrow windows centered at He #
584.34 Љ, O ### 599.60 Љ, O # 629.73 Љ and Mg ## 368.07 Љ.
As a complement to the temporal series, a 60 ## Ѕ240 ## context
image was obtained in three lines, He # 584.34 Љ, O # 629.73
Љ and Mg ## 368.07 Љ, alternating between the two studies.
The standard reduction was applied to the CDS data corн
recting for bias, flatнfield, cosmic rays, and instrumental e#ects
such as horizontal shifts due to the rotation of the scan mirror
and rotation and tilt in the spectrum due to the misalignment
between grating and detector, and grating and slit.
The small size and short lifetime of BPs added to the accuн
racy of the pointing, a few arcseconds, the small field of view
and the fact that the pointing has to be given a few hours before
the observing run, makes the choice of the target and the pointн
ing an important issue in the observations of these features. For
these reasons, `last minute pointing' (#3 hours before observн
ing) was required, using as a reference the latest EIT images.
The evolution, brightness and location of the BPs visible in the
previous hours were inspected to find the appropriate target,
i.e. a newly formed, bright and isolated (from active regions)
BP that would hopefully remain visible until the end of the obн
serving time.
2.2. Context images and magnetograms
As a context, we have EIT (Delaboudiniere et al. 1995) imн
ages in Fe ### 195 Љ and MDI (Scherrer et al. 1995) fullнdisk
magnetograms (1.96 ## /pixel), which cover with non regular caн
dences (6--90 minutes for the coronal imager and 1--90 minutes
for MDI) the lifetime of the two BPs. In addition, we have one
1024Ѕ1024 pixel TRACE Fe ##/# 171 Љ image showing the
fine structure of BP2 on November 8 at 00:34 UT.
In Fig. 1 the location of BP2 is highlighted with an arrow
in a fullнdisk EIT 195 Љ image, with a closeнup and a comparн

I. UgarteнUrra et al.: Signature of oscillations in coronal bright points 3
Fig. 2. Flux distribution for several EIT images of the same region,
with no bright point (dashed lines) and with its presence (solid lines).
The dotted line represents the threshold value chosen to define the BP.
ison with the TRACE image. It can be clearly seen that while
in EIT (2.6 ## /pixel) the BP appears as a di#use bright region
with a brighter central core, the TRACE (0.5 ## /pixel) image reн
solves it as an arcade of loops connecting opposite polarities.
The alignment with MDI was done by crossнcorrelating EIT
and TRACE images and matching the corrected TRACE coorн
dinates with MDI.
MDI magnetograms were corrected for the solar di#erenн
tial rotation (Howard et al. 1990) and geometrical projection
of the lineнofнsight magnetic flux (Chae et al. 2001; Hagenaar
2001). Standard routines were used in the reduction of EIT and
TRACE data (dark current subtraction, degridding, flatнfielding
and comic rays replacements).
3. Data analysis
3.1. BP identification
The identification of a BP is usually done establishing a thresн
hold value for its coronal emission. However, di#erent criteria
have been used in di#erent works. For example, Habbal et al.
(1990) used a factor of 2 higher flux than the surrounding quiet
region in Mg # 625 while Zhang et al. (2001) used a 3# (root
mean square) increase of the flux in an EIT 195 Љ quiet Sun
area. Here, we have chosen the threshold based on the flux disн
tribution (Brown et al. 2001). In Fig. 2 we plot the number of
pixels vs the 195 Љ flux for di#erent EIT images of the same
region with the BP emission still present (solid line) and when
it has already disappeared (dashed line). It shows a tail in the
distribution due only to the presence of the BP. We then choose
a threshold value (dotted line) over which the enhanced coronal
emission is considered as coming from a BP, which allows us
to define its lifetime.
3.2. CDS tracking
Compensation for solar rotation can be done in two ways with
CDS. Generally it is done by moving the rear legs that control
the pointing of the instrument, however, this is not the recomн
mended method when corrections of less than a few (3н5) arcн
seconds are required (Pike 1997). It is more suitable to make
use of the movement of the scan mirror, which reнpoints in
steps of 2 ## in the EнW direction. There is a small di#erence beн
tween the correction applied by the mirror, in fixed steps, and
the actual solar rotation. A crossнcorrelation of the two rasters
at the beginning and end of the temporal sequences plus a comн
parison with the solar Y profiles of the slit shows that the cuн
mulative e#ect after 120 minutes produces a drift of one pixel
in the EнW direction for our observing run.
3.3. Spectral analysis
It is well known in the CDS community that because of the
loss of contact with SOHO in 1998, the performance of the insн
trument changed after the recovery. Most of these changes are
well understood and can be dealt with, like the broadening and
asymmetry found in the postнrecovery spectral lines profiles of
the NIS detector (Del Zanna et al. 2001, see Fig. B.2). These
can be modeled with broadened Gaussian profiles (Thompson
1999). Nevertheless, extra care is needed in the subtraction of
the continuum because the broadening produces a blending of
the lines that hides the real continuum level in certain parts of
the spectrum. This problem, which a#ects NIS 1 to a higher
degree, can be avoided by extracting a wide spectral window,
as the one shown in Fig. 3, where the continuum can be easн
ily discriminated from the emission coming from blends and
fitted with a first and a zero order polynomials. For the secн
ond study we selected narrow windows, but the only one that is
registered in the NIS 1 detector is the Mg ## 368.07 Љ line and
there is no di#culty in determining the continuum from the red
wing. This line presents a blend with a Mg ### line at 367.68
Љ, which was extracted by assuming that both lines have the
same spectral width taking into account that they have a simiн
lar formation temperature. The other lines did not present any
complex blends. Line fitting was carried out using the standard
CDS package using broadened Gaussian profiles.
3.4. Time series
In order to get better signal to noise and taking into account
that the spatial resolution of CDS is larger than the pixel size
(Pauluhn et al. 1999), we summed the emission coming from
several pixels along the slit. We selected the central pixels of
the feature choosing in a quantitative approach those pixels
with an intensity 3 times the # level of the emission in the
raster, if available for that line, or along the slit otherwise. The
binning was done then along 21н22 pixels in Mg ##, 16 pixels in
O # and O ###, and 13н16 pixels in He #. In Fig. 4 we give an exн
ample of the procedure followed. On the left side, three rasters
for three of the above lines are shown with the pixel location
of the slit highlighted by a dashed line and a solid contour line
accounting for the 3# level. On the right hand side, the plots

4 I. UgarteнUrra et al.: Signature of oscillations in coronal bright points
Fig. 3. Sample spectrum of the wide spectral window extracted from NIS 1 in one of the studies. The dotted lines represent the two fits used to
subtract the continuum from the spectral lines. Some of the relevant ones are labeled.
Fig. 4. In the left panels are Mg ## 368, O # 629 and He # 584 rasters for the BP of 7 November 2002 (BP2). The slit location for the spectral
data is highlighted by a dashed line and a solid contour shows the 3# level. On the right hand panels, the plots correspond to the intensity
profile along the dashed line in the raster. The dotнdashed line is the threshold level and the dotted lines delimit the interval of integration.
correspond to the intensity profile along the dashed line on the
raster. The dotнdashed line is the threshold level and the dotted
lines delimit the interval of integration.
The appearance of BP2 in these rasters taken at the beн
ginning of the first time series is representative of the whole
dataset. The structure, although evolving and showing variaн
tions in intensity (to be discussed later), keeps approximately
the same size and configuration along the whole sequence in
the three lines, with the size in the transition region lines alн
ways slightly smaller than in the coronal line. This last stateн
ment fits the description given by Gallagher et al. (1998) for
several very bright network features in their study of the proн
perties of the quiet Sun EUV network.
The result of this binning are the temporal series to be disн
cussed later. The average cadence between observations is 94 s
for the run with a wide spectral window and 36 s. for the two
runs with narrow windows. The number of exposures (80 and
200) makes the total observing time #2 hours for each of them.
3.5. Wavelet analysis
The wavelet analysis is now a common technique used to anaн
lyze nonнstationary temporal signals. It has the advantage with
respect to the fast Fourier transform that it can localize freн
quencies in time, making it suitable for the study of temporal

I. UgarteнUrra et al.: Signature of oscillations in coronal bright points 5
Fig. 5. Evolution in time of the EIT flux, magnetic flux and EIT flux
per unit area (pixel) of the two consecutive BPs observed at the same
location. Dashed lines delimit their lifetime. The observed dataнpoints
are represented by plus symbols, triangles represent times with siн
multaneous EIT and MDI data for BP1 and stars for BP2. The black
dots are three points to be referred to in Fig. 6. The starting time is
November 4 2002, 00:00 UT.
series where di#erent frequencies and di#erent time localizaн
tions might be expected. A comparison of results obtained usн
ing both techniques can be seen for example in Lau & Weng
(1995) and Fludra (2001).
Software and definitions provided by Torrence & Compo
(1998) have been used in our study. The wavelet analysis conн
sists of a convolution of the time series with a wavelet funcн
tion resulting in a power spectrum, a two dimensional (time
and frequency) transform of the temporal series. We chose the
Morlet wavelet function. As in Fourier analysis, the highest
values of the power spectrum correspond to the relevant freн
quencies present in the signal, although in this case with the
time location also specified. This analysis su#ers from edge efн
fects at the limits of the time series. The region of the power
spectrum where these e#ects are important is called the cone
of influence (COI) and the results lying in this region should be
disregarded.
A crucial part of the analysis is to find the significance leн
vels in the power spectrum which tell us which are the real and
relevant frequencies of the signal. We used a Monte Carlo or
randomization method to estimate these levels (O'Shea et al.
2001). Randomization methods have the advantage that no asн
sumption about a noise model is needed. It just assumes that
if there is no periodicity in the time series then the intensity
values are independent of the time of occurrence: any order of
the intensity values in the time series would be as likely as any
other one. The test then consists of permuting n! times the temн
Fig. 6. Relation between the EIT Fe ### flux and the magnetic flux.
Same symbols as in Fig. 5 are used. The dashed line is the linear fit
to the BP1 points (triangles) excluding the three black dots, and the
dotнdashed line is the fit to the BP2 points (stars).
poral series, with n being the number of observations. Due to
computational and time constraints (O'Shea et al. 2001) we reн
stricted the calculations to a sample of n = 200 random permuн
tations. One then obtains the power spectrum for each of them
and compares the peak values with the ones obtained from the
data. The level of probability shown in our results is obtained
from the proportion of times that the random peaks are larger
than the peaks of the observed time series. We did that for the
two highest maximums (hereafter 1 st and 2 nd maximum) of the
power spectrum.
4. Results
4.1. Lifetime and evolution
According to the criteria of identification described earlier, the
BP studied with CDS (BP2) appeared in the EIT images on
November 7 at 04:48 UT and disappeared at 19:26 UT on the
8 th , which means that it had a lifetime of #38 h in the 195
Љ emission. From the magnetic field point of view the polarн
ities were present much earlier. In fact, another bright point
(BP1) was visible before this one at the same location. It apн
peared on November 4 at 17:36 UT with the emergence of a
positive polarity in an area of dominant negative polarity. After
several hours of interaction, the positive polarity canceled alн
most totally with the subsequent fade of the coronal emission
(November 6, between 20:48 UT and 22:24 UT), leaving a pair
of small positive fragments #10н13 Mm apart from the negн
ative one. After an interval of five hours, new negative flux
emerged at #6 Mm distance from the positive ones and BP2
appeared. During its lifetime, BP2 grew in area and intensity
lying always between the two polarities or covering the whole
bipolarity. The negative polarity being the dominant one with
several of the surrounding negative features joining in during
the process. The positive fragments, meanwhile, were graduн
ally and totally canceled by the negative ones until BP2 faded.
This dissipation occurred #3н4 hours before the full cancellaн
tion of the polarity, which took place some time between 22:24
UT and 00:00 UT on the 9 th .

6 I. UgarteнUrra et al.: Signature of oscillations in coronal bright points
Fig. 7. Flux variations for BP2 as seen in Mg ##, O #, O ### and He #. Start time is November 7 at 16:20 UT. The three sections correspond to
studies s26195, s26197 and s26199. Gaps in between were covered by context rasters. Dotted lines signal when the tracking correction took
place. The formation temperature of the spectral line is provided beside its identification.

I. UgarteнUrra et al.: Signature of oscillations in coronal bright points 7
In Fig. 5 the evolution in time of the EIT flux and magnetic
flux of the two bright points can be seen. The starting time is
November 4 at 00:00 UT. The EIT flux accounts for the inteн
grated emission of the pixels with flux over the identification
threshold. The total unsigned magnetic flux (# = # B l dS ) was
calculated integrating the line of sight magnetic field strength
(B l ) measured on the MDI magnetograms (300 s/exposure)
over a 60 ## square area below the coronal emission region usн
ing a cuto# value of ‘25 Gauss. The two top panels show the
important role that the magnetic field plays in the evolution
of BPs, as other authors have already shown (PreТs & Phillips
1999; Madjarska et al. 2003). However, taking into account that
the area evolves following a similar trend to the coronal emisн
sion, we should look at the evolution of flux per unit area (pixel)
at the bottom graph. We see that the behaviour of the EIT flux
is not only due to the increase/decrease of the BP area, but also
to the intrinsic evolution of the feature (Habbal & Withbroe
1981), which has a higher intrinsic emission when the magн
netic flux is at its maximum.
The dependence of the EIT flux on the total magnetic flux
can also be inferred from Fig. 6, where the symbols have the
same meaning as in Fig. 5. There is a linear dependence beн
tween these two quantities, excluding the three black dots.
They represent the points with maximum flux and area for BP1
and it seems that the high increase in intensity at these times is
not related to an increase of magnetic flux, which remains alн
most constant. It is interesting to note that the steepness of the
linear fits to the data points of both days (dashed line: BP1; dotн
dashed: BP2) is basically the same, which suggests that for siн
milar magnetic strengths BPs evolve similarly in intensity. The
fit is subjected to the uncertainty in the method used to calcuн
late the magnetic flux. However, introducing a time dependent
area of integration does not change the results and conclusions,
but marginally the steepness of the fit.
We have also checked the proposed dependence of the lifeн
time with the maximum area suggested by Golub et al. (1974):
A max = 2.5 Ѕ 10 7 # [km 2 ], where # is the lifetime in hours. The
predicted maximum area values are around 3 times larger than
the observed ones. BP1 has a lifetime of around 51 hours and
maximum area of around 4.8 Ѕ 10 8 km 2 , and BP2 38 hours and
3 Ѕ 10 8 km 2 , while the prediction gives 12 Ѕ 10 8 and 9.5 Ѕ 10 8
km 2 , respectively.
4.2. Time series
If we focus now on the CDS study, the temporal response of the
BP (BP2), as seen from di#erent lines corresponding to di#erн
ent temperatures of the solar atmosphere, is shown in Fig. 7.
The time series for each of the lines consist of three consecuн
tive sections: firstly the sequence run with longer exposure time
(81 s/exposure), followed by two sequences of 30 s/exposure.
The two gaps in between the three time series were covered by
context rasters. Other gaps are just missing data. The vertical
dotted lines account for the moments when the scan mirror was
moved for tracking. The plots show that the BP experiences
short time variations (order of minutes) in the flux. This manн
ifestation is clear for the two transition region lines (O # 629
Љ and O ### 599 Љ) and slightly smoother for the He # 584 line.
The coronal line (Mg ## 368 Љ) undergoes changes too, but on
a longer scale. The signal to noise is lower for this line, even
though we binned more pixels, and the error bars are larger.
However, it does not seem to hide these large fluctuations of the
flux seen in the cooler lines. These fluctuations are not ubiquiн
tous in time and can occur after a quiet period like the one seen
in the second section of the time series, 2 hours 27 minutes afн
ter the start of the series. The start time in the plot is November
7 at 16:20 UT.
The high activity seen in the transition region lines is highly
correlated at di#erent temperatures, as a glance to the two midн
dle panels shows. There is a one to one relation for the bursty
changes, and the general trend is certainly followed by the He #
line, reproducing even some of the peaks. Nevertheless, the
changes are less important with the maximum variation in the
amplitude being 40% at the beginning of the series, with an avн
erage of 10н30% from troughs to crests. For O #, the variations
are of the order of 30н60%. Within the temporal resolution of
our data, there is no noticeable time delay between the peaks at
di#erent temperatures.
4.3. Wavelet analysis results
It is clear from the present dataset that if we want to investigate
if there is any periodicity in these transient intensity changes,
we need to use a technique capable of detecting di#erent peн
riods or frequencies at di#erent times. As discussed before,
wavelet analysis can do this. We present now the analysis of
two BPs: first the one described in previous sections and next a
BP studied by Madjarska et al. (2003) with SUMER.
4.3.1. Present dataset
An example of the wavelet analysis carried out is shown in
Fig. 8. The top panel shows the time evolution of the intenн
sity in photonнevents. It is, in fact, a detrended intensity, result
of the subtraction to the original series of the linear fit to each
of the intervals between tracking corrections (dashed lines in
Fig. 7), as done by Fludra (2001). The power spectrum is shown
in the middle panel, with the level of probability corresponding
to the two higher maximums of the power for each time locaн
tion in the bottom panel. The 1 st maximum is represented by
thick dots in the power spectrum plot and by a solid line in
the probability one; the 2 nd by thin dots and by a dotted line.
The global wavelet spectrum, an average in time of the power
spectrum, lies on the rightнhand side.
A summary of the results is presented in Table 1 for each of
the three studies (s26195, s26197 and s26199) shown in Fig. 7.
An analysis was done for each of the three lines: Mg ## 368
Љ, O # 629 Љ and He # 584 Љ. The periods obtained for the
two highest maximums in the global wavelet spectrum are folн
lowed by the level of probability (1 - p), with p being the proн
portion in the number of permutations where the maximum of
the randomized power spectrum is higher than the maximum
of the real power spectrum. It gives an indication of how reн
liable the results are. Only periods with a level of probability

8 I. UgarteнUrra et al.: Signature of oscillations in coronal bright points
Fig. 8. BP2 wavelet results corresponding to study s26197 (middle series in Fig. 7) for O # 629.73 Љ (left) and He # 584.34 Љ (right). Top:
detrended intensity in photonнevents; center: wavelet power spectrum and global wavelet spectrum; bottom: level of probability for the two
highest peaks in the spectrum. 1 st maximums: thick dots and solid line; 2 nd maximums: thin dots and dotted line. The dashнdotted line correн
sponds to a probability level of 95%. Start time is November 7 2002 at 18:47 UT. The dashed line of the global wavelet plot represents the
maximum period outside the COI.
Fig. 9. Damped oscillation observed in the O # 629 time series, study
s26197. Overplotted in dotнdashed line there is an exponential damped
sine function with a period of 546 seconds.
higher than 0.95 (95%) are considered as possible candidates
for a periodic signal in this paper. These values are the result
of a comparison between the global wavelet spectra of real and
random series. Values of 1.0 just indicate that the probability
is between 99н100% (O'Shea et al. 2001). We will refer to the
power and probability plots for cases when time plays an imн
portant role. The relevant periods have been highlighted with
bold font. COI in the comments means that the maximum in
the power accounts for a signal which falls inside the cone of
influence, so it is disregarded. Other comments are referred as
numbers in the text.
The analysis of study s26195 in O # 629 Љ produces a peak
in the power at 1010 seconds with the probability over the 0.95
level (1), identified in the time series with several peaks with
a response time or duration close to the period. For He # 584,
where the changes are smoother, the analysis only detects an
initial peak inside the COI with a probability value of 0.93. It
is worth mentioning an isolated peak with a lifetime of 566
seconds experienced by O # at the beginning of the time series
(not given in Table 1).
Study s26197 shows a period of 1300 seconds with over
95% probability for the noisy Mg ## 368 Љ time series (2). An
inspection of the time dependent plots shows that it is due to

I. UgarteнUrra et al.: Signature of oscillations in coronal bright points 9
Table 1. BP2 wavelet results. Columns: study number, spectral line
used, maximum of the global wavelet power spectrum, period assoн
ciated with the maximum, probability (1 н p), and comments to the
results. COI: signal located inside the cone of influence; numbers: refн
erence to the text. The relevant periods are in bold font.
Study Line Max. Period [s] Prob. Comments
s26195 Mg ## 1 st 327 0.34
2 nd 1698 0.05
O # 1 st 1010 1.00 (1)
2 nd 2020 0.91
He # 1 st 1852 1.00 COI
2 nd 1010 0.93 COI
s26197 Mg ## 1 st 1300 0.99 (2)
2 nd 709 0.25
O # 1 st 596 1.00 (3)
2 nd 1003 0.98 COI
He # 1 st 919 1.00 (4)
2 nd 422 1.00 (5)
s26199 Mg ## 1 st 709 0.67
2 nd 1417 0.46
O # 1 st 546 1.00 (6)
2 nd 1002 1.00 (7)
He # 1 st 1838 1.00 COI
2 nd 596 1.00 (8)
the modulation seen in Fig. 7, some 3 hours into the observaн
tions. If we remove this low frequency with a 20нpoint running
average, there are no relevant periods remaining in the data.
The O # 629 Љ time series, however, show a very clear peak
at around 596 seconds (3), see Fig. 8, which is associated with
an interesting trend. Fig. 9 is a section of the series, just after
minute 60 from the start. There are three consecutive brightenн
ings with a similar response time ranging between 546 and 596
seconds, during a period where no tracking correction was necн
essary, followed after the correction by two more brightenings
of decreasing intensity with the peak of the wavelet inside the
COI. For a qualitative understanding of what clearly appears
as an oscillatory behaviour we have overplotted an exponential
damped sine function of amplitude
A(t) = A 0 sin # # # # # # #
2#t
P - # 0 # # # # # # # exp # # # # # # # -
t
t d
# # # # # # # (1)
with the initial amplitude A 0 = 4200 photonнevents, the period
P = 546 s. (obtained in the wavelet analysis), the phase # 0 =
7#/11 and the exponential decay time t d = 1800 s. The profile
reproduces well the intensity behaviour, suggesting that there
is a damped oscillation with a period that increases with the
damping. A strong brightening very similar to the initial ones
takes place just after (or even during) the weakest of the events.
Nothing can be said from the wavelet because its maximum is
inside the COI, but the duration, as seen in the light curve, is
from minimum to minimum #560 seconds, a typical response
time of other events.
In He # 584 Љ there is no sign of the damped oscillation,
however, the two first brightenings are also present with a maxн
imum in the wavelet at similar values for the period (Fig. 8
right). Preceding them there are high power spectrum values at
422 seconds that produce the second maximum in the global
wavelet (5), also seen for O # 629 Љ as a second maximum
(thin dots) on the left plot. At 20:03h UT, 76 minutes after the
starting point of the two plots in Fig. 8, one of these brightenн
ings experiences what seems like a flaring event with a 50%
increase in the O # 629 Љ flux of one of the datapoints, also
seen in Fig. 9. The same increase is found in the other transiн
tion region line, O ### 599, and a 15% increase in He # 584 Љ.
In the coronal line, although weaker, there is a jump of a 10%
just visible over the error bars. The length of the event is given
by the time resolution of the time series, i.e. 72 seconds. An
inspection of the unbinned data reveals that the event peaks in
a single spatial pixel spreading to the neighbours probably due
to the fact that the spatial resolution is larger than the pixel size
(Pauluhn et al. 1999). Finally, a 919 seconds period seems to
modulate the #500 seconds peaks at the center of the series in
He # during 3 complete cycles (4).
The last of the studies is s26199. As in the previous ones,
the coronal line time series does not reveal any relevant freн
quency over the noise. (6) O # 629 Љ shows again a peak at 546
seconds and the time dependent plot shows di#erent consecuн
tive brightenings with periods varying between 459 and 596
seconds at the start of the series. In He # 584 Љ, for the same
brightenings, we found the maximum in the wavelet power for
periods ranging from 459н649 seconds, with a maximum in the
global wavelet of 596 (8). The table shows also a peak at 1002
seconds in O # (7) that could modulate the center of the series.
Table 2. Wavelet results for the BP observed with SUMER on October
17, 1996. Columns: spectral line, maximum of the global wavelet
spectrum, period associated with the maximum and level of probaн
bility.
Line Maximum Period [s] Prob.
S ## 1 st 491 1.00
2 nd 901 1.00
4.3.2. SUMER BP
We have also inspected a BP observed on October 17 1996
with several SOHO instruments, where a detailed study was
done by Madjarska et al. (2003). We concentrate on the intenн
sity changes observed in the S ## 933.40 Љ line obtained with
the SUMER spectrometer and shown in Fig. 10 of that referн
ence. The observation was done in the sitнandнstare mode, i.e.
no tracking correction was applied. The time cadence of the seн
ries after a bin of 5 spectra (10 s/exposure) is 50 seconds. In that
figure the high intensity average values at the center of the time

10 I. UgarteнUrra et al.: Signature of oscillations in coronal bright points
Fig. 10. Wavelet results for the SUMER BP. See caption in Fig. 8 for
details. Units for the intensity are counts.
profile and the low ones at the wings are just a consequence
of the BP moving outside the slit, entering and leaving. This
trend was removed from the time series using a 30нpoint runн
ning average, and the result is plotted at the top of our Fig. 10.
The result of the wavelet analysis is shown in Table 2.
As explained in the cited paper, due to solar rotation the
SUMER slit (1 ## ) covered a new region of the Sun after #390
seconds, so that variations occurring during a period of 400н
500 seconds can be seen as a temporal change in a small scale
phenomenon. We have found four brightenings, after minute
40 from the start of the series, which have characteristic periн
ods in the range of 451н491 seconds, which means that these
are real temporal changes in the BP constituents. The oscillaн
tory nature of these intensity changes can be seen in Fig. 11.
A sinusoidal function with a period of 491 seconds, as derived
from the wavelet analysis, is overplotted on the detrended inн
tensity fluctuations. It is clear from the comparison of the two
curves that there is a regularity in the appearance of the brightн
enings. These 'short' brightenings are followed by longer ones
with a peak at 901 seconds.
4.4. Electron density
The temporal variation of the electron density (N e ) associated
with the intensity changes was also investigated. The determiн
Fig. 11. S ## detrended intensity fluctuations (solid line) corresponding
to the SUMER BP. Overplotted with a dashed line there is a 491 s.
period sinusoidal function.
nation of N e was done using the CHIANTI atomic database
(Dere et al. 1997; Young et al. 2003) and the ionization equiн
librium of Mazzotta et al. (1998). The identification of the
lines and possible blends was obtained using CHIANTI, the
list given in the quiet Sun EUV spectrum observed with CDS
(Brooks et al. 1999) and the list given by Del Zanna (1999).
Several density sensitive line ratios are present in the dataset
including ratios of Si ##, Si # and Fe ###. However, despite our
best e#orts, the postнrecovery broadening that a#ects the proн
files of the lines, mainly in detector NIS 1, does not make the
fitting procedure an easy task in the region where these lines are
located, namely in the range 340н347 Љ and 364 Љ in the case
of a Fe ### line (see Fig. 3). Although the fitting can be done
in certain cases, the uncertainties and assumptions involved do
not allow us to have su#cient confidence to present a tempoн
ral variation. The Si ## ratio 349.86/345.10 gives an average
value of #5 Ѕ 10 8 cm -3 consistent with the value presented
by Del Zanna et al. (2003) for a BP in a coronal hole using
a Si ## ratio. Line profiles from the second detector (NIS 2)
presented a better prospect and the O ## density sensitive line
ratio (625.8/608.4) was successfully determined. The 608.4 Љ
line is part of a set of several lines including the second order
Si ## 606.8 Љ and He ## 607.6 Љ lines, and a blend of a Mg # line
with an O ## one at 609.9 Љ. The 608.4 Љ O ## line was fitted
using the width obtained for O #, which is close in temperature
of formation. The 625.8 Љ line is located in the red wing of the
Mg # 625.0 Љ and even though is not a very strong line it can
be safely fitted.
As the BP is located in a coronal hole area, the background
emission is not very high. However, it is still important to do
a proper subtraction. For the 608.4 Љ line, we used a closeнby
region to the BP, where no network emission is seen and which
remains fairly constant. We avoided regions where bumps in
the emission could be due to network oscillations. The backн
ground emission of the 625.8 Љ was too weak to be fitted by a
Gaussian and could not be removed, so the N e values presented
here are an upper limit. If the background is not subtracted at
all, the general behaviour (shown in Fig. 12 and discussed beн
low) still remains the same, but the values are lower (as one
would expect) by on average 30%, which shows the importance

I. UgarteнUrra et al.: Signature of oscillations in coronal bright points 11
Fig. 12. Top figure: intensity changes along time for the two O ## lines
that form the electron density sensitive ratio 625.8/608.4. Bottom figн
ure: electron density changes along time. The start time is the same as
in Fig. 7. The dotted lines serve as reference for comparison.
of subtracting the background emission in order to get accurate
N e values.
Fig. 12 shows the temporal changes of the flux for the two
O ## lines. The start time is the same as in Fig. 7. The gap acн
counts for missing data and the vertical dotted lines serve just
as reference. The fact that the time changes are the same for
both lines is a good indication that the fitting and its assumpн
tions are reliable. The plot at the bottom shows the temporal
response of the electron density. Even though the errors, which
come from the propagation of the intensity errors through the
ratio, are large, from this plot we suggest that the intensity
changes are probably the result of electron density changes in
the transition region. Unfortunately, we do not have another raн
tio sensitive to electron density changes at this temperature to
confirm the result.
5. Discussion & Conclusions
The association of coronal bright points with bipolar magnetic
features was established in the first studies carried on BPs
(Krieger et al. 1971; Golub et al. 1977). These works were folн
lowed by many more concerning di#erent aspects of the magн
netic properties of BPs including statistical (Webb et al. 1993)
and individualized (Brown et al. 2001; Madjarska et al. 2003)
studies of the evolution of the photospheric magnetic features,
the analysis of the solar cycle dependence of the polarities
(Sattarov et al. 2002), its orientation with respect to the coronal
emission (Longcope et al. 2001) or the correlation with the Xн
ray and EUV emission (PreТs & Phillips 1999). We know from
the latter that the magnetic flux determined from the photoн
spheric magnetic features is well time correlated with the Fe ###
coronal emission and the Xнray flux, in the way shown in the
top and middle panels of our Fig. 5. Our contribution extends
the sample to two more BPs, adding the result that not only
does the coronal emission rise due to an increase in emission
area (which increases too), but also because the BP emits more
per unit area as the magnetic flux becomes stronger. Therefore,
there is more intrinsic emission when there is more magnetic
flux involved. In fact, we find there is a one to one relation beн
tween the magnetic flux in the bipolarity and the EIT Fe ###
coronal emission for the two BPs, both in the growing and deн
caying phase. The exception is the brightest moments of the
first BP when the high increase in emission is not accompanied
by a comparable increase in the magnetic flux. The reason for
that is not clear at the moment and more BPs have to be studied
to see if there is a saturation limit in the magnetic strength.
These results give further support to the picture that shows a
BP whose emission is a result of the interaction of the magnetic
polarities, with most likely magnetic reconnection involved in
the process. The converging flux model proposed by Priest
et al. (1994) describes how the approach of two opposite polarн
ities creates a Xнpoint that rises into the corona and produces
a Xнray bright point (Xнray emission) by coronal reconnection
(see also Parnell et al. 1994; Longcope 1998). The model proн
poses a three phase evolution: preinteraction (approach), interн
action and finally cancellation. Several observational questions
are raised concerning the evolution of the features. For examн
ple, the timing between coronal emission and magnetic cancelн
lation. In this case, BP1 fades (Fe ### 195 Љ emission below the
threshold) when, after the interaction and almost full cancellaн
tion of the positive polarity, the remaining fractions of opposite
polarity are #10 Mm apart. BP2 appears when new positive
flux emerges at a distance of #6 Mm and finally disappears
#3н4 hours before the full cancellation of the positive polarity.
This is in agreement with Harvey et al. (1999) who observed
several examples of coronal emission disappearance before the
cancellation of magnetic network elements and suggested that
magnetic flux is submerging at most of these cancellation sites.
It would be interesting to check with new observations the timн
ing of disappearance of the emission at di#erent temperatures.
Another key aspect of the evolution of BPs is their high
variability in EUV lines and Xнrays. In a set of several papers
and looking at spectroheliograms obtained from Skylab experн
iments, Habbal and collaborators (Habbal & Withbroe 1981;

12 I. UgarteнUrra et al.: Signature of oscillations in coronal bright points
Habbal et al. 1990; Habbal & Grace 1991) analyzed the intenн
sity fluctuations of several chromospheric (Ly #, C ##), transiн
tion region (C ###, O ##, O ##) and coronal (Mg #) lines. They
reported significant temporal fluctuations between scans, findн
ing shortнterm variations of the order of #5 minutes, but also
more gradual ones (20н30 minutes). These fluctuations were
correlated between transition region lines, but not always with
variations at other temperatures. The variability was most enн
hanced at 10 5 K and no characteristic periodicities were found.
They also pointed out that both quiet Sun and coronal hole
bright points behave similarly, concluding that the BP properн
ties are independent of the structure of the overlying largeнscale
magnetic field. The main limitation of the studies was the temн
poral resolution of their data, the 5.5 minutes needed to obtain
the spectroheliogram. Our observations with an improved temн
poral resolution (94 and 36 s) have confirmed some of these
results. We have found variations in the flux of He # 583.33 Љ
and O # 629.73 Љ with a characteristic response time ranging
between 420н650 seconds. Many of these events appear in a
random fashion and sometimes after periods of quietness. The
strongest variability is in the transition region lines, well correн
lated between them, with no counterpart in the changes of the
Mg ## 368.07 Љ emission. The coronal emission for this BP
seems to evolve following more gradual changes.
The interpretation of the behaviour of He # 583.33 Љ is more
di#cult. It is under debate how the helium lines are formed in
the Sun and several mechanisms have been proposed over the
years. Which process is the dominant could determine where
to expect helium to emit. If photoionization from coronal raн
diation followed by recombination is dominant, it could be
formed in the upper chromosphere, while if that mechanism
is not the dominant one, the formation could take place at the
lower transition region (Andretta et al. 2003, and references
therein). This last scenario would explain why the He # series
follow so closely the trend given by the transition region lines
and why it does not follow the coronal one. However, to discuss
the He lines formation is not the purpose of this paper.
Finally, what it is important to stand out from our obserн
vations is the oscillatory behaviour present in the two bright
points to which we applied the wavelet analysis. In the SUMER
BP, looking at the S ## 933.40 Љ intensity fluctuations we have
found a peak in the wavelet spectrum over a 95% significance
level at 491 s. during at least four cycles. The sitнandнstare
mode was used and new plasma was seen under the slit evн
ery #400 seconds. The oscillatory pattern could be then exн
plained by fourнfive small structures (# 1 ## ) evenly spaced,
which moved under the slit producing the periodic intensity
fluctuations as they pass. However, it seems to us much more
likely to have a 4н5 ## structure crossing the slit while it experiн
ences oscillations in its emission. Similarly we have found osн
cillations in the emission of the O # 629.73 Љ line of the CDS
bright point. The wavelet power spectrum has its peak at 546
s. In this case the oscillation decreases in amplitude producing
a damped oscillatory profile with a change in the period in the
last few cycles.
The wavelet analysis also provides periods of 900н1000
seconds for both SUMER and CDS suggesting that there could
be a longer modulation component. In the case of the SUMER
observations this modulation is only two cycles long. Since
new plasma is seen under the slit every 400 seconds it could
be that this is just a response to the morphology of the bright
point. These uncertainties together with the realization that the
fine structure of the bright point can only be `resolved' with
the highest spatial resolution images, as shown from the comн
parison of EIT and TRACE images, suggests that coordinated
studies of high spatial and temporal resolution images should
go hand in hand with high cadence spectroscopic studies, if
one wishes to understand better the nature of the oscillations.
This issue will be adressed in a forthcoming paper as well as
an extended study looking for oscillations at higher temperaн
tures, which have already been found in other coronal features
(Aschwanden et al. 1999; O'Shea et al. 2001, and references
therein).
Acknowledgements. Research at Armagh Observatory is grantнaided
by the N. Ireland Dept. of Culture, Arts and Leisure. This work
was supported in part by PPARC grant PPA/G/S/1999/00055 and
PPA/V/S/1999/00628. CDS, EIT, MDI and SUMER are instruments
onboard SOHO. SOHO is a project of international cooperation beн
tween ESA and NASA. CHIANTI is a collaborative project involving
the NRL (USA), RAL (UK), and the Universities of Florence (Italy)
and Cambridge (UK).
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