JuanCastillo29/space-mission-intelligence-data huggingface dataset:
Fitted Jovian spectra excluding 0.94 keV feature (04 October 2011)
10−4
10−4
keV (Photons cm−2 s−1 keV−1)
keV (Photons cm−2 s−1 keV−1)
10−5
10−5
10−6
10−7
(data−model)/error
(data−model)/error
−2
1 0.5 2
1 0.5 2 −2
Energy (keV)
Energy (keV)
Bestfit Chandra spectra on 02 Oct, 2011
Bestfit Chandra spectra on 04 Oct, 2011
10−4
10−4
keV (Photons cm−2 s−1 keV−1)
keV (Photons cm−2 s−1 keV−1)
10−5
10−5
10−6
10−7
(data−model)/error
(data−model)/error
−2
−2
1 0.5 2
1 0.5 2
Energy (keV)
Energy (keV)

[Figure: Figure 4. Top panels show Chandra/ACIS fitted X-ray spectra obtained from the northern polar region (black crosses) and southern polar region (red circles) of Jupiter on 2nd October 2011 (top left panel) and 4th October 2011 (top right panel), with a combination of a Bremsstrahlung and Gaussian emission line near 0.73 keV. A significant excess is observed near one keV in residual plots during both dates. The excess can be mitigated and the best-fit obtained by adding another Gaussian line near 0.94 keV, as shown in the bottom-left and bottom-right panels for 2nd and 4th October 2011, respectively. Model components used to fit all spectra are shown by thin dotted lines.]
exchange interaction of hydrogen ions. Therefore, both com-
ponents of Lyα lines may originate from different geograph-
ical locations of Jupiter: the broader one may come from
the aurora, where UV-emitting gas has strong dynamics,
while the narrow component may originate from the disc,
dominated by resonantly scattered or thermally broadened
hydrogen emission. Throughout our work, we treat both line
fluxes separately. In addition to Lyα, best-fit spectra require
several Lorentzian functions depending on the number of
emission lines. As an example, the best-fit model of the ob-
= 164.67/155, with the best-fit spectra, model components,
and residuals shown in the bottom panel of Figure 3 in 1150-
1450 ˚A . We derive integrated fluxes for both these emission
lines and the continuum (1370-1420 ˚A) using the cflux con-
volution model in XSpec, which provides fluxes over chosen
wavelength ranges along with their 1σ errors. Each of the 51
spectra is fitted separately to obtain flux measurements at
every pointing. The resulting fluxes are then corrected for
Jupiter’s distance to ensure consistency across all observa-
tions.
continuum, thereby highlighting systematic deviations as-
sociated with distinct spectral features. Ratio plots clearly
show that there are two Feii line complexes near 1600 ˚A dur-
ing solar minima, which are missing from the other spec-
trum, raising the possibility that the UV spectral lines be-
tween the Solar minima and Solar maxima can be different.
To track such changes over time, we extract continuum and
line light curves by fitting each spectrum in XSpec using a
mix of narrow and broad Lorentzian components to capture
the emission features and the underlying continuum. As seen
in the top panel of Figure 3, a single Lorentzian component
does not fit the Lyα emission line and a strong line resid-
ual is observed around 1218 ˚A . Hence, we fit two Lorentzian
components to describe the particular line. While 1213 ± 1.6
˚A line is broad in all spectra (best fit values vary between
0.032-0.045 eV), 1218 ± 1.8 ˚A
line is consistently narrow
(best fit values vary between 0.011-0.023 eV). Branduardi-
Raymont et al. (2007) reported broad oxygen emission lines
in XMM-Newton spectra, indicating ion velocities of order
dergoing charge exchange in Jupiter’s magnetosphere. In our Tomer et al.
[Figure: Figure 5. Top panel shows 1, 2 and 3σ integrated contours from the marginal probability distribution of Ne8+ emission line energy near 0.94 keV versus line normalisation as obtained from the bestfit X-ray spectral parameters shown in the bottom left panel of Figure 4. Bottom panel shows Chandra/ACIS image extracted in the 0.8–1.2 keV energy range from the observations conducted on 2nd October, 2011 (ObsID 12315).]
dra/ACIS were modelled using the spectral fitting package
XSPEC v12. Owing to the small number of spectral bins
(10–15), the Kolmogorov–Smirnov (KS) statistic was pre-
Due to the red light leakage in soft X-ray bands and op-
tical loading effect, the soft X-ray flux below 0.5 keV is
contaminated beyond corrections in Chandra/ACIS spectra
of Jupiter (Elsner et al. 2005). However, contamination is
negligible in Chandra/HRC spectra. Therefore, we have ex-
cluded ACIS data analysis below 0.5 keV. 0.5-8 keV spectra
obtained with Chandra/ACIS were initially modelled using
a bremsstrahlung continuum (bremss) following earlier stud-
ies (Branduardi-Raymont et al. 2007). Significantly peaked
residuals were observed, motivating the inclusion of the
Gaussian emission line in the model. For example, for Ob-
sID 12315, fitting with only the bremsstrahlung model gives Jupiter’s North pole 0.04
Time (minutes; 0 min = 22 September, 2011 07:40 AM)

[Figure: Figure 6. The Solar X-ray activities, including a CME event, are shown along with the subsequent Chandra/ACIS flaring lightcurves of Jupiter on 2nd October 2011 (red) and 4th October 2011 (blue). For clarity, the zoomed plot of the X-ray flare as observed on 4th October is shown in the inset.]

[Table: Table 4. 0.5-8 keV Best-fit parameters for the Chandra/ACIS spectra of Jupiter’s northern and southern auroral region, with the width (σ) of the first Gaussian line fixed at 0.06 keV, and the listed fluxes include distance corrections. detailed description of obs1 and obs2 is provided in Table 2.
Component Parameter(unit) obs1 obs2 gaussian LineE(keV) 0.74+0.01 −0.01
0.72+0.02 −0.02
norm(N) 7.37+1.31 −1.32
7.28+1.60 −1.71
norm(S) 3.60+1.00 −0.98
2.04+0.97 −1.33
gaussian LineE(keV) 0.98+0.02 −0.02
0.96+0.05 −0.07
sigma(keV) 0.07+0.02 −0.02
0.14+0.04 −0.04
norm(N) 3.74+1.18 −0.93
3.67+2.23 −1.53
norm(S) 2.34+0.73 −0.63
3.44+1.77 −1.23
bremss kT(keV) 1.28(f) 1.01+0.57 −0.28
norm(N) 13.68+1.87 −1.95
9.54+5.90 −4.34
norm(S) 5.27+1.45 −1.49
4.24+3.04 −1.94 flux0.5−10(N) (10−14) 4.34+0.19 −0.19
3.12+0.21 −0.21 flux0.5−10(S) (10−14) 1.94+0.18 −0.18
1.55+0.14 −0.14] KS test (log) -3.91 -4.52
solar activity influences Jupiter’s auroral X-ray flux. Exam-
ination of the light curves reveals a notable intensification
in the auroral emission by a factor of ∼3 following the solar
event. The temporal correlation between the increase in so-
lar flux and the corresponding rise in Jupiter’s auroral X-ray
flux is further illustrated in Figure 6, highlighting the poten-
tial impact of solar activity on Jovian auroral processes. In
particular, enhancements in the 0.5–10 keV Chandra/ACIS
light curve appear approximately 8–10 days after a strong,
recurring CME event associated with an M-class solar flare
recorded between 24 and 26 September 2011. This temporal
analysis complements the spectral results by providing in-
sight into the short-term variability of the auroral region as
described in the next section.
dra/ACIS observations are analysed to study the temporal
behaviour of Jupiter’s northern auroral X-ray emission. The
datasets corresponding to Observation IDs 12315 and 12316,
taken on 2nd and 4th October 2011, respectively, were anal-
ysed using identical data reduction and light-curve extrac-
tion procedures. The solar X-ray activity surrounding the
Chandra/ACIS X-ray flux of Jupiter on 2nd and 4th October
2011 is shown in Figure 6. The background-subtracted light
curve for the northern auroral region for Observation ID
12316 is shown in the inset of Figure 6. These observations
were selected due to limited data availability, and they were
obtained 8 and 10 days after a major CME was reported,
HRC lightcurve of Jupiter’s NPR on 10−April−2014
HRC lightcurve of Jupiter’s NPR on 15−April−2014
Total counts = 723
Total counts = 531
Count/sec
Count/sec
0 0.02 0.04 0.06
0 0.02 0.04 0.06
104 2×104 3×104 4×104
104 2×104 3×104 4×104
Time (sec)
Time (sec)

[Figure: Figure 7. Chandra/HRC X-ray flares and CME. The top left panel shows the Chandra/HRC image of Jupiter in 0.5-10 keV as observed on 20th April, 2014. The source and background regions used for analysis are marked in the image. The top right panel shows the 0.5-10 keV background-subtracted Chandra/HRC lightcurves for Jupiter’s northern auroral region as observed on 10th April, 2014 (shown in red circles) and 15th April, 2014 (shown in blue squares), along with the corresponding solar X-ray flux (in black solid line) starting from 1st April, 2014. A CME was reported and shown as observed on 03 April, 2014. For the clear visibility of X-ray flares, zoomed versions of both light curves are shown in the bottom left and bottom right panels, respectively.]
the
temporal
variability
of
Jupiter’s auroral X-ray emission, we analysed a total of
29 Chandra/HRC observations spanning between 08 April
2014 and 15 September 2021(see Table 3). Background-
subtracted light curves were extracted for all observations,
and both flare and non-flare periods were examined to de-
termine whether enhancements in Jupiter’s X-ray flux occur
regularly or are linked to elevated solar activity. Two repre-
sentative light curves, corresponding to 10 April 2014 and 15
April 2014, are shown in the bottom left and bottom right
panels of Figure 7, respectively. On the top right panel of
Figure 7, the Chandra/HRC observations are shown for 10
and 15 April, along with the corresponding solar X-ray flux
starting from 1st April, 2014, to provide the broader solar
activity context leading up to and during the Jupiter ob-
servations. The enhanced X-ray flares can be observed from
zoomed-in lightcurves shown in the bottom panels of Figure
7. Since the hot spot is located near Jupiter’s north pole
and is fixed in longitude, it is visible for only about half of
each Jovian rotation. As a result, one would expect the ob-
served count rate to increase when the hot spot rotates into
and 1300-1400 ˚A continuum emission for the year 1978-1986
(21st solar cycle) in the left and right panels of Figure 8
and the same are shown in the left and right panels of Fig-
ure 9 for the year 1986-1996 (22nd solar cycle), which high-
lights the long-term variation of Jupiter’s flux. We compare
these fluctuations with simultaneous sunspot number data
obtained from the silso database. During the 1986–1996 pe-
riod, X-ray flux data from the Sun, measured by the GOESJupiter Ly  UV Light Curves from the Pearson correlation coefficients between the Lyα flux and the solar flux, as well as between the continuum flux and the solar X-ray flux, to investigate their correlation.
For
Lyα at 1218 ˚A, we obtain correlation coefficient r = 0.678
with a p-value of 0.002 and for Lyα at 1213 ˚A r = 0.551
with p = 0.018. While for the continuum emission r = 0.89
tions indicate that increases in solar flux are reflected in
Jupiter’s ultraviolet emission. In several cases, particularly
at the time of solar minima, we have also observed two sig-
nificant peaks around 1575 ˚A and 1608 ˚A. These features
are likely due to stacking of multiple, unresolved Feii lines
(at 1575.81, 1575.42, 1575.02 ˚A and at 1607.98, 1608.29,
1608.45, 1608.53, 1609.03, 1609.03, 1609.1 ˚A). The spectral
fitting of these two cases is shown in Figure 10. We have
detected Feii line complexes in 12 IUE spectra, spanning
between 1982 and 1994. Interestingly, we have detected the
Siii line at 1812 ˚A during observation ID SWP25688 with a
The line is shown in Figure 11. We calculated fluxes for both
Feii lines during 12 observations. We compare Feii line flux
variabilities with Lyα lines and show them in Figure 12. The
fluxes of both lines are fairly constant during rising, peak
or descending phases of Lyα line variabilities (which follow
X-ray flux variabilities during solar cycle) as shown in Fig-
ure 12. Such a non-correlation indicates that Feii may not
originate from solar interactions with Jupiter’s atmosphere;
rather, activities of Io may influence such origin, which we
explore in the discussion section. Another possibility is the
presence of iron in solar wind or magnetosphere, which may
be photoionized by solar radiation. The line flux correspond-
ing to Siii is denoted as a magenta star in Figure 12.

1978 to 1986 and Solar Activity
UV Flux (1218 Å)
4.5
UV Flux (1213 Å)
4.0
Sunspot number
3.5
Sunspot Number
UV Ly  Flux
3.0
2.5
2.0
(10 11 erg cm 2 s 1)
1.5
1.0
1978 1980 1982 1984 1986
Time (Year) Jupiter UV Continuum Light Curves from
1978 to 1986 and Solar Activity
UV Continuum Flux
1.8
Sunspot number
1.6
1.4
Sunspot Number
UV CONT Flux
1.2
1.0
0.8
(10 11 erg cm 2 s 1)
0.6
0.4
1978 1980 1982 1984 1986 0.2
Time (Year) Figure 8. UV Lyα and continuum (1300-1400 ˚A) flux variation of IUE/SWP data from 1978 to 1986. UV flux is plotted alongside sunspot number data from the same period.
with the simultaneous solar activity parameters like X-ray
fluxes and Sunspot numbers, we found that UV Lyα line and
continuum fluxes from Jupiter’s atmosphere evolve signifi-
cantly: during solar maxima to minima or vice versa, the UV
Lyα line flux changes by a factor of ∼3.5 (as observed from
top panels of Figure 8) while the same during UV continuum
changes by a factor of ∼4-5 (as observed from bottom panels
of Figure 8). Figure 8 and Figure 9 show moderate to strong
correlation (Pearson’s correlation coefficient 0.68-0.91) with
the evolution of solar X-ray fluxes as well as sunspot num-
Jupiter Ly  UV Light Curves from
1986 to 1994 and Solar Activity
FeII
10−12
ergs cm−2 s−1 Å−1
4.0
UV Flux (1218 Å)
3.5
UV Flux (1213 Å)
5×10−13
3.0
Solar X-ray Flux Sunspot no.
Solar X-ray Flux
2.5
UV Ly  Flux
2×10−13
2.0
(data−model)/error
1.5
(10 6 W m 2)
1.0
(10 11 erg cm 2 s 1)
−2
0.5
1560 1580 1600 1620 1640 Wavelength (Å)
0.0
FeII
1986 1988 1990 1992 1994 1996 1
10−12
ergs cm−2 s−1 Å−1
Time (Year) Jupiter UV Continuum Light Curves from
5×10−13
1986 to 1994 and Solar Activity
4.0
Continuum Flux
2.25
3.5
Solar X-ray Flux Sunspot no.
2×10−13
(data−model)/error
2.00
3.0
Solar X-ray Flux
UV CONT Flux
1.75
2.5
1.50
−2
2.0
1.25
1.5
1560 1580 1600 1620 1640
(10 6 W m 2)
Wavelength (Å)
1.0
1.00
(10 11 erg cm 2 s 1)

[Figure: Figure 10. Spectral fitting in the range of 1550-1650 ˚A. Top: The best-fit (χ2/d.o.f. = 50.61/50) UV spectrum is shown for the Observation ID SWP31444. Bottom: Best-fit (χ2/d.o.f. = 35.54/51) spectrum for the Observation ID SWP20739. The best-fit model (red) consists of continuum and emission line features, and the residual of the fitting is shown in the bottom panel.]

0.5
0.75
0.0
1986 1988 1990 1992 1994 1996 0.50
Time (Year)

[Figure: Figure 9. UV Lyα and continuum (1300-1400 ˚A) flux variation of IUE/SWP data from 1986 to 1994. UV flux is plotted alongside solar X-ray flux and sunspot number data from the same period. The sunspot numbers are normalised with respect to solar X-ray flux data.]

FeII SiII
ergs cm−2 s−1 Å−1
10−13
10−14
cycles; Jupiter’s polar Lyman-α (1215.67 ˚A) emission and
far-UV continuum are highly sensitive to variations in solar
extreme ultraviolet (EUV) and UV (UV) irradiance. This
sensitivity arises because the upper atmosphere and exo-
sphere of Jupiter contain a significant population of neutral
hydrogen atoms, which act as resonant scatterers for incom-
ing solar Lyα photons. The scattering process dominates the
baseline Lyα brightness across the polar regions and is mod-
ulated directly by changes in the solar UV flux (Clarke et
al. 1980; Gladstone, Waite, & Lewis 1998) Solar EUV and
X-ray photons play a crucial role in photochemistry in the
Jovian upper atmosphere. High-energy UV and X-ray pho-
tons ionise and dissociate molecular hydrogen into atomic
hydrogen, populating the exosphere and thus increasing the
column density of atomic hydrogen for resonance scatter-
ing. Therefore, during times of peak solar activity marked
by elevated sunspot numbers and observed EUV and X-ray
fluxes, photodissociation becomes maximally efficient, and
the density of scattering H atoms becomes very high. These
scattering H atoms increase Lyα emissions and far-UV con-
tinuum in the auroral areas. The continuum emission arises

10−15
(data−model)/error
2 0 −2
1780 1800 1820 Wavelength (Å)

[Figure: Figure 11. The best-fit (χ2/d.o.f. = 41.42/37) UV spectrum in 1760-1840 ˚A range is shown for the Observation ID SWP25688. The best-fit model (red) consists of emission line features, and the residual of the fitting is shown in the bottom panel.]
related with the solar cycle were studied by Clarke et al.
(1980) and Gerard et al. (1993) using IUE and Copernicus
observations, leading to the discovery that polar Lyα fluxes
increase during maxima and diminish during minima of the
solar cycles. Results from our work are also being confirmed Jovian UV-X-ray 13

FeII line at 1608 Å
Ly  line at 1218 Å
FeII line at 1575 Å
Ly  line at 1213 Å
Si II line
UV Flux
(10 11 erg cm 2 s 1)
1982 1984 1986 1988 1990 1992 1994
Time (Year)

[Figure: Figure 12. UV Lyα and Feii line flux variation of IUE/SWP data from 1981 to 1994. Fluxes for both the Feii lines are fairly constant as compared to the significant variation for the Lyα lines. The magenta star denotes the observation time of Siii line detection and its flux value.]

cluding iron, into Jupiter’s magnetosphere (Bagenal & Delamere 2011). These ions become entrained in the corotating plasma and are transported along magnetic field lines toward the auroral regions, where precipitation leads to collisional excitation of metal ions and emission in the UV. The flux of such Feii lines is therefore strongly modulated by magnetospheric dynamics, for example, rotational modulation of the magnetosphere, episodic plasma injections, interchange instability events, or Io-related plasma variations- rather than by the long-term solar cycle.
by earlier findings by Bhardwaj & Gladstone (2000) and Gustin et al. (2016).
On non-correlation of solar activities with ionised UV
lines

Our results show that fluxes due to ionised UV lines like Feii at 1608 ˚A and 1575 ˚A (Figure 10) show no correlation with the solar activities studied over 12 years, as observed from Figure 12. Ionised metal lines, such as Feii at 1575 ˚A and 1608 ˚A, in Jupiter’s UV spectra exhibit a fundamentally different origin than the Lyα and far-UV continuum. Unlike Hydrogen lines, which are dominated by solardriven resonant scattering, these metal lines cannot be efficiently excited by solar UV photons alone because the crosssections for resonance scattering of Feii are small, and the solar flux at the relevant wavelengths is weak (G´eRard et al. 2006; G´erard et al. 2019). Instead, observations and modelling suggest that Feii and other heavy ion emissions are primarily generated by interactions of precipitating magnetospheric ions with the upper atmosphere (Yamauchi et al. 2024). These ions can be accelerated within Jupiter’s magnetosphere via processes such as centrifugal interchange instability, or magnetotail reconnection events, producing collisional excitation of Fe+ ions in the atmosphere of Jupiter (Grodent et al. 2003; Kimura et al. 2015). A particularly important source of such ions is the Io plasma torus. Io, with its intense volcanic activity, injects large quantities of neutral and ionised sulphur, oxygen, and trace metals, in-
On the origin of SiII line

The detection of Siii (1812 ˚A) in Jupiter’s spectra during solar minima, shown in Figure 11 and Figure 12 is particularly intriguing because it appears only when the solar-driven UV background is low, suggesting that the line is either inherently weak or masked by stronger solar-scattered emissions during periods of high solar activity. This behaviour indicates that the observed Siii line is likely dominated by local or magnetospheric processes rather than by direct solar excitation. One plausible explanation is enhanced visibility of weak, locally generated emissions during solar minima. When the solar UV flux is low, the background Lyα and far-UV continuum decrease, reducing the competing signal from resonance scattering. This allows intrinsically faint auroral emissions, such as Siii produced by particle precipitation or low-density plasma processes, to become detectable (G´eRard et al. 2006; G´erard et al. 2019). A second posTomer et al.
the magnetotail. These processes can temporarily increase
the flux of magnetospheric metal ions along field lines into
the polar atmosphere, producing a transient enhancement
in Siii emission (Bagenal & Delamere 2011; Artemyev et al.
2020). To differentiate among these possibilities, time-tagged
spectral observations combined with in-situ magnetospheric
measurements (e.g., from JUNO or future missions) or si-
multaneous remote sensing of auroral morphology would be
critical.

Sunspot Number HRC Count Rate
0.12
flares observed with Chandra/ACIS (Figure 6) and two flares
from Chandra/HRC (Figure 7), 7-15 days after major CMEs
were detected. Due to the scarcity of Jupiter’s X-ray obser-
vations after CMEs, the beginning and the end of such flares
cannot be determined. Using same Chandra/ACIS observa-
tions in 2011, Dunn et al. (2016) found that due to an in-
terplanetary coronal mass ejection (ICME), Jupiter’s X-ray
aurora intensified by nearly a factor of eight, accompanied by
rapid changes in spatial, spectral, and temporal behaviour,
including two hotspot periodicities from distinct ions: 26
min from sulphur ions and 12 min from a carbon/sulphur
and oxygen. Their analysis showed that sulphur-dominated
while mixed oxygen–sulphur–carbon emissions map to more
distant or open field lines, suggesting a possible solar-wind
contribution. Together with coincident radio bursts, they
observed that both the auroral enhancement and the be-
haviour of the polar X-ray hot spot were driven by ICME-
induced magnetospheric compression and/or pulsed dayside
reconnection. It has been proposed that enhanced solar-wind
disturbances associated with CMEs can compress Jupiter’s
magnetosphere and trigger dayside magnetopause reconnec-
tion, leading to plasma injections and magnetospheric re-
configuration (Nichols, Cowley, & McComas 2006). To test
whether observed X-ray flares in the lightcurve and highly-
ionised emission line in the spectra (bottom panels of Figure
7) can be induced by CME events, we calculate the closeby
coronal mass arrival time on the Jovian atmosphere.
On 2011 September 24, multiple coronal mass ejections
(CMEs) were detected with notably high velocities§. The
first major event occurred around 09:48 UT, associated with
an X1.9 flare at approximately 09:21 UT, followed by a sec-
ond eruption at about 12:48 UT, connected to an M7.1 flare
at 12:33 UT. During both the X-ray CME launch and the
subsequent Chandra observations (after a week), the Sun-
Jupiter-Chandra geometry was highly favorable, with Chan-
dra viewing Jupiter at a very small phase angle (observer
almost fully illuminated view of Jupiter’s northern auroral
region and enabling a direct investigation of the Jovian re-
sponse to solar activity. The CME speed measured is roughly
SOHO/LASCO catalogue¶ with the kinetic energy of 3 ×
[Figure: Figure 13. Top panel shows long-term (2014-2021) comparison of X-ray count rate (in log scale; after distance correction) obtained from each Chandra/HRC observations with sunspot numbers. The distance-corrected count rate values are given in Table 3. The middle and bottom panels show zoomed-in segments of the HRC light curve during solar maxima (2014) and solar minima (2019), respectively (shown by red circles), and corresponding solar X-ray flux values (shown in blue solid line). Neither long-term sunspot numbers nor short-term X-ray flux correlates with Chandra/HRC X-ray count rate from Jupiter.]
welling, which can selectively populate the excited states
electrons or heavy ions, auroral potential drops, or localised
reconnection events may energise ions in the polar upper at-
mosphere, leading to collisional excitation of Siii (Grodent
et al. 2003; Kimura et al. 2015). Such events may be sporadic
and independent of the long-term solar cycle, explaining why
the line appears only intermittently. A third mechanism is
magnetospheric reconfiguration, such as plasma injections,