UPDATE ON TITAN’S
NIGHT-SIDE AIRGLOW
P. LAVVAS, R.A. WEST
UVIS Team meeting - Jan 2014
1
indicat
dayglo
(bottom
is dom
and wh
than th
[8] T
so we
1000–1
In Fig
synthe
then sm
lution
or syst
multip
same c
FUV (
(0.5 #
with th
multip
the day
“Bright” disk emission (I/F~10-7)
Weak limb emission
West et al. 2012
Figure 1. (top) UVIS Titan FUV airglow data on 7 May
2009 between 8:27 and 10:50 UT while in Saturn’s shadow.
The composite fit to the data is overplotted in red and the
altitudes over which the data is averaged are indicated. (top
middle) Same as Figure 1 (top) but showing the component
2
of the fit from the N2 LBH bands Radiances and uncertain-
How can we explain these observations?
2.2. T
[9] T
sure tim
UV (31
coverin
et al., 2
era as
system
ratio in
from t
Figure
[10]
In Figu
3
APPROACH
The observed night-glow is due to the
excitation of N2 by non-solar energy sources.
METHOD
1. Identify energy sources
II. Calculate N2 state population & local emission rates
III. Simulate radiation transfer of emitted photons
through atmosphere and instrument
(for limb and nadir geometry)
IV. Compare with observations
4
I. POTENTIAL ENERGY SOURCES
1200
1000
Altitude (km)
800
600
Magnetospheric electrons
~1000 km
Magnetospheric ions (O+)
~800 km
Magnetospheric protons
~800 km
~500 km for silicates
~700-800 for water ice
Meteoroids
400
200
(depending on B)
Chemi-luminescence
~200 km for C2H2
Cosmic Rays
~65 km
5
I. POTENTIAL ENERGY SOURCES
Gronoff et al. 2009, 2011
O+
off et al.: Ionization processes in the atmosphere of Titan. I.
Fig. 7. Ionization from mono-energetic protons, with an isotropic distri- Fig. 9. Ionization due to mono-energetic atomic oxygen ions O+ , with
bution computed with planetocosmics. The higher in altitude the peak, Fig.
an3.isotropic
distribution
computed with
Thea higher
Ionization
due to EUV/XUV
solar planetocosmics.
flux computed for
solar inFig. 5. Comparison of electron impact ionization computed with
the less in
energetic
the protons.
altitude
theofpeak,
the less
energetic
the ions.
transTitan and planetocosmics.
zenith
angle
40◦ with
a solar
flux given
by F10.7 = 80.
G. Gronoff et al.: Ionization processes
the atmosphere
of Titan. I.
Protons
Electrons
+
Fig. 13. Ionization
due todue
oxygen
(O(O
)+precipitation
with
energy
Fig. 13. Ionization
to oxygen
) precipitation with
thethe
lowlow
energy
Fig. 10. Ionization due to precipitating protons for the T5 (active) con-Fig. 6. Electron impact ionization computed with the Cassini T5’s flyby
Fig.
4. Ionization from mono-energetic electrons, with an isotropic dis- electron flux conditions with transTitan. A vertical
protons
spectrum
considered
as
oxygen.
G. Gronoff
et al.:
processes in the a
spectrum
as oxygen.
magnetic flux
line Ionization
is
Fig. protons
8. Ionization
due considered
to mono-energetic
isotropically distributed pro- ditions, computed with the coupled model.
for quiet
conditions.
pitating
protons
for quiet conditions.
tribution computed with planetocosmics. The higher in altitude the
tons, with energy near the cutoff, computed with planetocosmics. The peak, the less energetic the electrons. The aim of this figure is to show considered.
higher in altitude the peak, the less energetic the protons.
the peak altitude with respect to the energy. In the following, the elec−3 −1 TransTitan, which is more accurate
tron
is computed
through
aninfluence
amplitude
of 0.3 cm
s . The T5 conditions (Fig. 10) corshown in Fig. 5. The electron production due to the electron preand
is
able
to
discriminate
between
species.
responds to active conditionstheofionized
protons;
the peak is close to
reduced to fit the energy detection range of the Cassini probe.
While in Fig. 8, we plot the influence of the protons in the
10 MeV–1 GeV range, in such energies we are at the limit between the protons accelerated in the magnetosphere of Saturn
and the cosmic rays. In these figures, the secondary electrons
(produced by proton-ionization) are taken into account as an energy deposition at the altitude where they are produced, whereas
in the coupled code, the electrons are introduced into the kinetic
part.
The interesting point is that the protons with energy above
1 MeV have an ion production peak below 500 km. In Fig. 11,
the ion production from proton precipitations is presented for
quiet conditions. The production peak is close to 750 km with
cipitated flux peaks at 900 km below the EUV–XUV peak. The
850 km with an amplitude of 6 cm−3 s−1 . In these two plots, a
amplitude of production is about 5 cm−3 s−1 . We compared this
secondary layer shows up between 400 km and 600 km, with ancomputation with the results of the full transport equation solufrom
1 keV slightly
up to 100
keV (not
in the figamplitude
smaller
than all
0.1profiles
cm−3 s−1shown
. It originates
in the
ure).
The
production
was
computed
on
an
average
based
on tion, shown in the same figure, and the values compare very well.
precipitations of protons with energies above 1 MeV
as shown
more
than7.3000 runs on each angle. The same numerical mode The accuracy is better above about 1100 km with the transport
in Fig.
code, because the production at high altitude stem from to low
was used
planetocosmics
study. In
The
The with
influence
can be seenall
in through
Fig. 9 forthis
oxygen.
thiselecfigure,
trons typically ionize above 550 km. The 100 keV electrons pro- energy fluxes that are not taken into account in planetocosmics.
the secondary electrons are taken into account as an energy deduce a maximum at 550 km and the 1 keV electrons at 900 km. One of the reasons is that the 35 eV value per creation of ion
position at the altitude where they are produced, whereas in the
The amplitudes of production on this plot cannot be interpreted electron pair goes wrong below these energies.
coupled code, the electrons are introduced into the kinetic part.
directly. They correspond to a production for an input flux of
equalcm
energy,
−2 −1 the oxygen ions have an influence at higher
1 At
electron
s on top of the ionosphere (and are therefore 4.3. Ionization from precipitating protons
of precipitation
create aninion
layer
−2 −1
inaltitude.
unit of This
(cm−3kind
s−1 )/(cm
s ). To cannot
be interpreted
term
of below
600
km.
In
Fig.
12,
we
present
the
ion
production
due
production, this value has to be multiplied by the input spec- toIn Fig. 7, we present the ionization from precipitating protons
precipitating
ions. Following
et al. (2008),
trum
from Fig. oxygen
1 and integrated
over theCravens
energy range.
This is wewith different characteristic energies. The set in energy has been
gen (O+ ) precipitation (the T5 protons
en here).
Fig. 14. Ionization due to mono-energetic cosmic ray protons, with an
isotropic distribution computed with Planetocosmics. The higher in altitude the peak is, the less energetic the cosmic ray are.
tudy that the whole T5 precipitation
itation
T5 toprotons
he
peak (the
is close
900 km with an
r quiet conditions, (Fig. 13) the peak
−1
production peak
is at
65 km with an amplitude
15 cm−3 swith
. an
Fig.of14. Ionization
due to
mono-energetic
cosmic rayof protons,
tude of 0.4 cm−3 s−1 . In the case
the peak,
the production
due to the cascade
down toin alhere is no secondary layer below
the Below
isotropic
distribution
computed
with Planetocosmics.
The higher
ground
is less
clearly
visible. the
In Fig.
17, ray
we are.
compare the reis, the
energetic
cosmic
whole T5 precipitation titude thethepeak
sult for low solar activity (350 MV) with highest solar activity
se to 900 km with an
Cosmic Rays
Fig. 15. Ionization due to galactic cosmic rays computed with the coupled model.
6
Fig. 17. In
I. POTENTIAL ENERGY SOURCES
Production rates for
different N2 states
7
NOTES ON SIMULATION OF NIGHT AIRGLOW
II. AIRGLOW MODEL
3
Table 1. Radiative transitions included in the model.
Transition
Band Name
Max A[s−1 ] hν(0,0) [nm]
> 5000 transitions
emission
1 Σ+
Triplet A 3 Σ+
→
X
Vegard-Kaplan
1.950E-01
201
u
g
3
3
+
B Πg ↔ A Σu
First Positive
1.110E+05
1046.9Cartwright, 1978
3
3
W ∆u ↔ B Πg
Wu-Benesch
5.400E+03
135830.8
Gilmore et al. 1992
"
3
−
3
B Σu ↔ B Πg
2.290E+04
1527.9
Campbell et al. 2010
3
3
C Πu → B Πg
Second Positive
1.310E+07
337
3 +
E 3 Σ+
Herman-Kaplan
3.100E+03
217.3
g → A Σu
3
+
3
E Σg → B Πg
1.561E+02*
3
E 3 Σ+
7.635E+02*
g ↔ C Πu
3
D 3 Σ+
Fourth Positive
3.779E+07*
u → B Πg
Collisions
3
+
3
+
D Σu → E Σg
Quenching,
Singlet w 1 ∆u → X 1 Σ+
1.250E+08
139.5
g
w 1 ∆u ↔ α 1 Πg
MCF2
4.770E+03
3640
Energy pooling,
1
1
+
α Πg → X Σg Lyman-Birge-Hopfield 5.160E+03
145
V-V,
V-T,
α 1 Πg ↔ α" 1 Σ−
MCF1
2.620E+03
8252.5
u
1 +
Inter/intra-system crossing
α" 1 Σ−
Fifth Positive
u → X Σg
b 1 Πu → X 1 Σ+
Birge-Hopfield
g
1 +
c"4 1 Σ+
Carroll-Yoshino
u → X Σg
1
c"4 1 Σ+
u → α Πg
Doublet A 2 Πu ↔ X 2 Σ+
Meinel
5.850E+04
1109.2
g
2 +
B 2 Σ+
First Negative
1.140E+07
391.2
u ↔ X Σg
stribution of the states. Therefore, scaling the energy spectrum of the photoelectrons might result
me discrepancies from the actual energy spectrum of the produced electrons from cosmic rays or
ospheric electrons.
lines
Steady state solution for each state/level:
Production
Loss
3. Airglow simulation
simulation for the generation of airglow takes into account the production rates for the different
and calculates
the
rates
of de-excitation
due
toP
radiative transitions
and collisions
with
other
(1
−
Q
)(P
+
+
P
)
=
L
+
L
PD
direct
cascade
collisions
cascade
collisions
es. The calculations include the states and radiative transitions presented in Table 1. For the
8
nal transitions we have included intersystem and intrasystem collisional
processes for the single
II. N2 states population
9
II. LOCAL EMISSION RATES
10
III. Radiation Transfer (ISS)
DISK (below 300km): Optically thick, use DISORT
code with emission rates as source function and
consider extinction by aerosols (Tomasko et al.
2008) and CH4 (Karkoschka 1998).
300 km
Limb (above 300 km): Optically thin, integrate emission
rates along line of sight in a spherical shell geometry.
10.000
1.000
Ae
ros
ols
Photons (arbitrary units)
100.000
CH4
Intensity at TOA for a spectrally flat source located
between 50 and 70 km altitude, demonstrating clear
signatures of aerosols and methane extinction.
0.100
0.010
0.001
200
400
600
800
Wavelength (nm)
1000
11
with ISS measurements of scattered sunlight to character
aerosols within the layer. At first, the lack of this layer in s
the occultations appears to challenge the idea that the de
III. Radiation Transfer (UVIS)
Integrate emission rates along line of sight in a spherical shell geometry and consider attenuation by
hydrocarbons and aerosols with(a)
abundances from photochemical model (Lavvas et al. 2008) and UVIS
observations (Koskinen et al. 2011).
ty profile of C2H4 retrieved
h a vertical wavelength of
ean density was introduced
Opacity spectrum for
emitted photons at UVIS/
FUV range
f wavelength and imtions. The T41 I occul-
Emitted photons probe reasonably
well the whole FUV range.
(b)
Observed opacity
spectrum at UVIS/FUV
range during T53
occultation.
12
IV. COMPARISON WITH OBSERVATIONS
Night
Day
10 of 17
Complete Band Emissions
Night profiles = Day profiles /10
(roughly)
13
Stevens et al. 2011
IV. COMPARISON WITH OBSERVATIONS (UVIS)
L18204
WEST ET AL.: TITAN AIRGLO
ind
day
(bo
is
and
tha
[
so
10
In
syn
the
lut
or
mu
sam
FU
(0.
wi
mu
the
Figure 5. Vertical profiles of limb emissions for different bands.
2. Comparison between simulated
al., 2012]. All values are in Rayleighs.
emissions
and
UVIS
observations
Band emissions (R)
Model
Observation
Total FUV
LBH 10.35 8.86
7.2±5.1
VK
39.40 1.11
3.9±2.6
Atomic 0.24 0.23
4.4±1.4
West et al. 2012
Band
14
Figure 1. (top) UVIS Titan FUV airglow data on 7 May
2009 between 8:27 and 10:50 UT while in Saturn’s shadow.
The composite fit to the data is overplotted in red and the
altitudes over which the data is averaged are indicated. (top
middle) Same as Figure 1 (top) but showing the component
of the fit from the N2 LBH bands Radiances and uncertainties from the fit over the wavelength region in Figure 1
(top) are indicated. (bottom middle) Same as Figure 1 (top
2.2
[
sur
UV
cov
et
era
sys
rat
fro
Fig
[
In
atm
IV. COMPARISON WITH OBSERVATIONS (ISS)
CL1,CL2
CL1,BL1
IR2,CL2
CL1,VIO
Disk at μ=0.5
Limb Average (300-1000 km)
CL1,CL2 CL1,VIO CL1,BL1
IR2,CL2
CL1,CL2 CL1,VIO CL1,BL1
IR2,CL2
Observed
11
0.2
1.1
0.5
71.7
0.7
4.5
6.5
Model
4.35
0.17
0.42
0.39
0.1
0.005
0.014
0.004
Ratio
2.53
1.18
2.62
1.28
717
140
321
1625
15
OTHER POTENTIAL CONTRIBUTIONS
I. Star Light
OPTICAL EBL. I. RESULTS
57
band filters and photometers, as were used by Dube et al.
(1977, 1979). Finally, IRAS has provided maps of the thermal emission from dust at high Galactic latitudes. We have
used the IRAS maps to select a line of sight for these observations that has a low column density of Galactic dust in
order to minimize the DGL contribution caused by dustscattered starlight and also to estimate the low-level DGL
that cannot be avoided.
Our measurement of the EBL utilizes three independent
data sets. Two of these are from HST: (1) images from the
Wide Field Planetary Camera 2 (WFPC2) using the F300W,
F555W, and F814W filters, each roughly 1000 Å wide with
central wavelengths of 3000, 5500, and 8000 Å, respectively,
and (2) low-resolution spectra (300 Å per resolution element) from the Faint Object Spectrograph (FOS) covering
3900–7000 Å. The FOS data were taken in parallel observing mode with the WFPC2 observations. While flux calibration of WFPC2 images and FOS spectra achieve roughly
the same accuracy for point-source observations, the
increase in spatial resolution, a 104 times larger field of view,
lower instrumental background, and absolute surface
brightness calibration achievable with WFPC2 make it betFig. 1.—Relative surface brightnesses of foreground sources, upper
ter suited than FOS to an absolute surface brightness mealimits on the EBL23 (see x 1), and lower limits based on the integrated flux
of the EBL.
the-1FOS observations
5 surement
-2 sNonetheless,
-1 nm-1 sr
AB mag)light
in the HDF
(Williams et~10
al.
from resolved galaxies
(V555 > 23 Star
Direct
source:
photons
cm
do provide a second, independent measurement of the total
1996). The spectral shape and mean flux of zodiacal light and of DGL are
-1 nm
-1 of terrestrial airof -2
thesnight
sky,-1
also
shown at the levels we detect
in this work. The
airglowsource:
spectrum is taken
Zodiacal
light
~104 background
photonsflux
cm
srfree
from Broadfoot & Kendall (1968) and is scaled to the flux level we observe
glow and extinction, but with greater spectral resolution
CL1/CL2
DN~10
(observed
= images.
71.7)The third data set consists of longat 3800–5100 Å (see x 9). The effective bandpasses for our HST observathan the WFPC2
tions are indicated at the bottom of the plot.
slit spectrophotometry of a region of ‘‘ blank ’’ sky within
the WFPC2 field of view. These data were obtained at the
‘‘ blank screens,’’ spatially isolating all foreground contribu- 16 2.5 m du Pont telescope at the Las Campanas Observatory
tions from the background. This pioneering work produced
(LCO) using the Boller & Chivens spectrograph simultane-
OTHER POTENTIAL CONTRIBUTIONS
II. Chemi-luminescence
C2H2 ~300 ns
C4H2 ~100 ms
T. SHIRAI et al.
Electron Collisio
GRAPHS. Cross Section vs
See page 155 for Explanati
III. CH4 emission
Only from dissociation fragments
CH (420-440 nm) but too weak
17
CONCLUSIONS
I. Strongest emissions (ISS) in the upper atmosphere come
from the Vegard-Kaplan, 1st positive and Meinel bands, while
for the lower atmosphere 1st negative and 2nd positive
dominate. For UVIS/FUV Vegard-Kaplan, LBH and atomic N
emissions dominate.
II. Simulated limb emissions are consistent with observations
and are dominated by nominal magnetospheric energy input.
III. Disk observations are much higher than the simulated
emissions. Contributions by star light is 10x smaller that
observed signal.
18
BACK-UP SLIDES
19
!"##$%$&
Figure 1.
View from above the north pole of Saturn showing the positions of Saturn,
Titan and Cassini during the eclipse event described in the text. Saturn’s rotation and
plasma flow are in the counter-clockwise direction. The sun is off the right edge of the
20 rings.
frame and Saturn’s shadow is shown on the
Raw image