# First Detection of [SI] in Near-IR JWST Observations of Io in Eclipse, and Comparison With SO Emissions, Evolving Volcanic Eruptions, and Prior UV HST-STIS [SI] Emissions

Imke de Pater, Darrell F Strobel, Ashley Gerard Davies, Joachim Saur, Lorenz Roth, Katherine de Kleer, Emmanuel Lellouch, Zachariah Milby, Carl Schmidt, Thierry Fouchet, et al.

► To cite this version:

Imke de Pater, Darrell F Strobel, Ashley Gerard Davies, Joachim Saur, Lorenz Roth, et al. First Detection of [SI] in Near-IR JWST Observations of Io in Eclipse, and Comparison With SO Emissions, Evolving Volcanic Eruptions, and Prior UV HST-STIS [SI] Emissions. Journal of Geophysical Research. Planets, 2025, 130 (11), 10.1029/2024JE008850. hal-05341109

HAL Id: hal-05341109

https://hal.sorbonne-universite.fr/hal-05341109v1

Submitted on 31 Oct 2025

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Distributed under a Creative Commons Attribution 4.0 International License
JGR Planets

RESEARCH ARTICLE

10.1029/2024JE008850

Special Collection:

Solar System Science from
JWST

Key Points:

* The first detection of 1-micron sulfur lines on Io, compared with visible and UV emissions, reveals a system that is stable over decades

* Although simultaneously obtained SO emissions are produced by volcanism, the S emissions suggest excitation through electron impact

* Volcanic thermal emission is detected from the Kanehekili region and Loki Patera. SO, but no S, emissions are evident above Kanehekili

Correspondence to:

I. de Pater,
imke@berkeley.edu

Citation:

de Pater, I., Strobel, D. F., Davies, A. G., Saur, J., Roth, L., de Kleer, K., et al. (2025). First detection of [SI] in near-IR JWST observations of Io in eclipse, and comparison with SO emissions, evolving volcanic eruptions, and prior UV HST-STIS [SI] emissions. *Journal of Geophysical Research: Planets*, 130, e2024JE008850. https://doi.org/10.1029/2024JE008850

Received 20 NOV 2024
Accepted 14 AUG 2025

Author Contributions:

**Conceptualization:** Imke de Pater

**Data curation:** Imke de Pater.

Lorenz Roth, Michael H. Wong,

Julie Renaud-Kim, Carolyn Jordan

**Formal analysis:** Imke de Pater, Darrell
F. Strobel, Ashley Gerard Davies

**Funding acquisition:** Imke de Pater

**Investigation:** Imke de Pater, Darrell

F. Strobel, Ashley Gerard Davies

**Methodology:** Imke de Pater, Darrell

F. Strobel, Ashley Gerard Davies,

Emmanuel Lellouch

**Project administration:** Imke de Pater

**Resources:** Imke de Pater

© 2025. The Author(s).
This is an open access article under the
terms of the Creative Commons
Attribution License, which permits use,
distribution and reproduction in any
medium, provided the original work is
properly cited.

<figure><img src="image_2.png" /></figure>

First Detection of [SI] in Near-IR JWST Observations of Io
in Eclipse, and Comparison With SO Emissions, Evolving
Volcanic Eruptions, and Prior UV HST-STIS [SI] Emissions

Imke de Pater¹,² ID, Darrell F. Strobel³ ID, Ashley Gerard Davies⁴ ID, Joachim Saur⁵ ID, Lorenz Roth⁶ ID, Katherine de Kleer⁷ ID, Emmanuel Lellouch⁸ ID, Zachariah Milby⁷, Carl Schmidt⁹ ID, Thierry Fouchet⁸ ID, Michael H. Wong¹ ID, Leigh N. Fletcher¹⁰ ID, Jake Harkett¹⁰ ID, Michael Roman¹⁰,¹¹ ID, Richard J. Cartwright¹² ID, Jason E. Perry¹³, Julie Renaud-Kim¹⁴, Carolyn Jordan¹⁴, Stephan Schlegel⁵ ID, David A. Williams¹⁵ ID, and David M. Nelson¹⁵ ID

¹Department of Astronomy, University of California, Berkeley, CA, USA, ²Department of Earth & Planetary Science, University of California, Berkeley, CA, USA, ³Departments of Earth & Planetary Sciences and Physics & Astronomy, Johns Hopkins University, Baltimore, MD, USA, ⁴Jet Propulsion Laboratory, California Institute of Technology, Pasadena, CA, USA, ⁵Institute of Geophysics and Meteorology, University of Cologne, Koeln, Germany, ⁶Space and Plasma Physics, KTH Royal Institute of Technology, Stockholm, Sweden, ⁷Division of Geological and Planetary Sciences, California Institute of Technology, Pasadena, CA, USA, ⁸LIRA, Observatoire de Paris, Université PSL, CNRS, Sorbonne Université, Université Paris Cité, Meudon, France, ⁹Department of Astronomy, Boston University, Boston, MA, USA, ¹⁰School of Physics and Astronomy, University of Leicester, Leicester, UK, ¹¹Facultad de Ingeniería y Ciencias, Universidad Adolfo Ibáñez, Santiago, Chile, ¹²Johns Hopkins University Applied Physics Laboratory, Laurel, MD, USA, ¹³University of Arizona, Tucson, AZ, USA, ¹⁴W. M. Keck Observatory, Kamuela, HI, USA, ¹⁵Arizona State University, Tempe, AZ, USA

**Abstract** We observed Io with the James Webb Space Telescope (JWST) NIRSpec/Integral Field Unit (1.0–5.3 µm, *R* ~ 2700) in August 2023 while the satellite was in eclipse. Thermal emission from Kanehekili Fluctus is consistent with the cooling of lava flows after a vigorous eruption in November 2022. At Loki Patera, after a new brightening event was detected in November 2022, the lava lake was in a quiescent state, as expected from previous analyses. We mapped the SO emission band at 1.707 µm, and detected, for the first time, [SI] emission lines at 1.082 and 1.131 µm. The SO emissions are concentrated above Kanehekili Fluctus, and in two regions in the northern hemisphere. The disk-averaged brightness is 14.5 kR. The emissions are sourced from SO molecules ejected from 1,500 to 1,700 K vents in an excited state, with a typical SO column density above the northern hemisphere of ~10¹⁵ cm⁻². Sulfur emissions are distributed homogeneously across a band in the northern hemisphere. The disk-averaged total brightness is 5.6 kR, versus 9.65 kR in the north. The emissions are produced through direct electron impact excitation by ~4 eV electrons in the torus (density 2,500 cm⁻³) penetrating the atmosphere, and require the atmosphere to be hot (~1,700 K) to populate the upper levels before excitation. The sulfur column density over the northern hemisphere is 8 × 10¹⁴ cm⁻². These same parameters can explain recent 0.7725-µm observations, as well as the 147.9-nm multiplet emissions observed with HST-STIS >20 years earlier. This suggests a quite stable system over decades-long timescales.

**Plain Language Summary** We observed Io with the James Webb Space Telescope (JWST) between 1.0 and 5.3 µm in August 2023. The satellite was in Jupiter's shadow, so band/line emissions would not be obscured by Io's reflected sunlight. Volcanic thermal emission was detected from Loki Patera and Kanehekili Fluctus. The main changes in the shape of the thermal emission spectra since JWST observed Io in November 2022 were consistent with the continued cooling of emplaced lava flows in the Kanehekili Fluctus region, and the crust that had formed on the surface of the lava lake in Loki Patera. Images of Io in the SO 1.707 µm emission band show a concentration above Kanehekili Fluctus and in two regions in the northern hemisphere. The emissions are sourced from SO molecules ejected from volcanic vents. We further detected, for the first time, sulfur line emissions at 1.08 and 1.13 µm. These emissions are distributed homogeneously across a band in Io's northern hemisphere. They are mainly produced through excitation by electrons from the plasma torus, penetrating Io's atmosphere. This process predicts 0.7725 µm line emission in agreement with recent Keck observations, as well as UV emissions at the same brightness as observed with Hubble Space Telescope's in 2000. This shows that the system (plasma torus and atmosphere) is quite stable.
**Software:** Imke de Pater, Darrell F. Strobel, Ashley Gerard Davies

**Validation:** Imke de Pater, Darrell F. Strobel, Ashley Gerard Davies

**Visualization:** Imke de Pater, Darrell F. Strobel, Ashley Gerard Davies, Joachim Saur, Lorenz Roth, Zachariah Milby, Jason E. Perry, Stephan Schlegel

**Writing – original draft:** Imke de Pater, Darrell F. Strobel, Ashley Gerard Davies

**Writing – review & editing:**
Imke de Pater, Darrell F. Strobel, Ashley Gerard Davies, Joachim Saur, Lorenz Roth, Katherine de Kleer, Emmanuel Lellouch, Zachariah Milby, Carl Schmidt, Thierry Fouchet, Leigh N. Fletcher, Jake Harkett, Michael Roman, Richard J. Cartwright, Jason E. Perry, David A. Williams, David M. Nelson

# 1. Introduction

Io, the most volcanically active body in our Solar System continues to surprise us. Strong tidal heating induced by the 4:2:1 Laplace orbital resonance between Io, Europa, and Ganymede (Peale et al., 1979) manifests through ongoing, energetic volcanic eruptions. These are visible through thermal radiation from volcanic hot spots at near- and mid-infrared wavelengths, and through plumes of gas and dust in the UV and visible wavelength range (Davies, 2007; de Kleer & Rathbun, 2023; de Pater, Goldstein, & Lellouch, 2023; Lellouch et al., 2007). The composition of Io's atmosphere and some gaseous plumes has been characterized via spectroscopic observations, starting with the detection of SO₂ gas above Loki in *Voyager I*IRIS spectra (Pearl et al., 1979), followed by a detection of Io's global SO₂ atmosphere (Lellouch et al., 1990). Since that time, Io's atmosphere and plumes have been observed across the electromagnetic spectrum, from the UV through mm-wavelengths including imaging, disk-averaged and disk-resolved spectroscopy (de Pater, Goldstein, & Lellouch, 2023).

The abundance and spatial distribution of SO₂ gas has been derived from observations in the UV, IR, and mm range. For example, by observing the Ly-α line in the UV, SO₂ has been mapped since it absorbs the surface-reflected solar Ly-α radiation (Feaga et al., 2009; Giono & Roth, 2021). SO₂ has been observed directly in absorption against Io's background at mid-IR (Spencer et al., 2005; Tsang et al., 2012) and near-IR wavelengths (Lellouch et al., 2015), and in emission at mm wavelengths (Lellouch et al., 1990; Moullet et al., 2008). Observations in the UV have revealed, for example, the presence of S₂ in Pele's plume (Spencer et al., 2000), and equatorial auroral spots in the [SI] 190 nm and [OI] 136 nm transitions above Io's limb (Feaga et al., 2002; Roesler et al., 1999). At mm-wavelengths, both KCl and NaCl gases have been detected in localized spots above the surface (de Pater, Luszcz-Cook et al., 2020; Redwing et al., 2022). In the near-IR, the (forbidden) 1.707 µm SO emission complex has been imaged while the satellite was in eclipse. Intriguingly, there appeared to be no clear connection between the SO spatial distribution and the location of hot spots (de Pater et al., 2007; de Pater, de Kleer, & Ádámkovics, 2020), and the wing of the SO emission band near 1.69 µm was highly suggestive of non-LTE effects (i.e., gas that did not get thermodynamically equilibrated) (Bernath & Bittner, 2020; de Kleer, de Pater, & Ádámkovics, 2019; de Pater et al., 2002; de Pater, de Kleer, & Ádámkovics, 2020).

To further investigate the relation between the SO 1.707 µm emission band complex and volcanic hot spots, we observed Io with the James Webb Space Telescope (JWST) on 15 November 2022 as part of our Early Release Science (JWST-ERS) program #1373. We observed the satellite while in eclipse since the bands/lines cannot be detected against bright Io when in sunlight. We used the Integral Field Unit (IFU) on NIRSpec with the G235H/F170LP and G395H/F290LP settings, that is, obtained spectra between 1.65 and 5.3 µm at a spectral resolution R ~ 2,700. At that time there was an energetic eruption at Kanehekili Fluctus and a new brightening event had just started at Loki Patera (de Pater, Lellouch, et al., 2023) (see also Section 3.2). Maximum temperatures over both volcanoes were measured at ~1,200 K over areas as small as 0.06 km² over Loki Patera and 0.25 km² over Kanehekili Fluctus.

In our 2022 JWST data, we detected the 1.707 µm forbidden SO emission band complex at the location of Kanehekili Fluctus, which was the first time that this emission had been clearly seen above an active volcano. This suggests that the origin of such emissions is likely ejection of SO molecules in an excited state directly from the vent, as initially proposed by de Pater et al. (2002) based upon disk-integrated spectra. Low emission levels were also seen over large fractions of Io's surface, perhaps originating from stealth volcanism. This type of volcanism arises from extremely high-entropy eruptions (Kieffer, 1982) originating from a reservoir of superheated SO₂ vapor generated from the thermal exchange between hot silicate magma and a liquid reservoir of SO₂ about 1.5 km below the surface, at pressures of ~40 bar and temperatures of ~1,400 K. Since such plumes would consist of essentially pure gas, that is, without dust or condensates, they cannot be detected in reflected sunlight, and hence were not seen by spacecraft. Johnson et al. (1995) suggested this type might be widespread on Io, and was inferred to be the source of much of the JWST and other telescopic SO observed emissions (de Pater, de Kleer, & Ádámkovics, 2020; de Pater, Lellouch, et al., 2023). Unfortunately, the 2022 JWST observations at the shortest wavelengths, 1-1.8 µm with G140H/100LP, failed. These data should have yielded a robust determination of the highest temperature components present at Io's volcanic hot spots, and also would have provided much higher quality SO spectra, due to much longer integration times planned in the G140H grism compared to G235H. The experiment was repeated in August 2023. Although no particularly energetic eruptions were seen (preventing determination, or at least constraint, of lava eruption temperature), this new data set revealed new and surprising results, which are presented in this paper.
<figure><img src="image_3.png" /><figcaption>**Figure 1.** Geometry of Io in eclipse observations as seen from James Webb Space Telescope. On 23 August, Io entered eclipse at 17:19 UT, and the first integration started at 17:44 UT. Io's movement with respect to Jupiter is shown for each of the four dithers. On 4 August, Io entered eclipse at 6:05 UT, and the first integration commenced at 06:40 UT (Table 1). Io's motion during the observations is indicated, shifted upwards by 4".</figcaption></figure>

The current paper is organized as follows: In Section 2 we discuss the observations and data reduction techniques. We then split the paper into 4 parts: Section 3: Io's thermal hot spot emissions, Section 4: the SO emission band, Section 5: the S emission lines, and Section 6: non-detections of O₂, S₂O, and S₂. Conclusions are provided in Section 7.

## 2. JWST Observations and Data Reduction

A repeat of the 15 November 2022 JWST-ERS #1373 (de Pater, Lellouch, et al., 2023) observations was conducted on UT 04 and 23 August 2023. As before, observations were obtained with the IFU of JWST's near-infrared spectrograph NIRSpec. On 4 August, the same experiment as on 15 November 2022 was carried out, so that spectral data cubes were obtained in two settings: (a) G395H/F290LP, covering 2.87–5.27 µm, with an effective exposure time of 42.9 s (or total duration of 64.4 s) in each of 4 dithers, and (b) G235H/F170LP, covering 1.66–3.17 µm, with an effective exposure time of 85.9 s (or total duration of 128.8 s) in each of 4 dithers. On 23 August, a third spectral data cube was obtained: (c) G140H/F100LP, covering 0.97–1.89 µm, with an effective exposure time of 583.6 s (or total duration of 641.9 s) in each of 4 dithers. Together these grisms cover a wavelength range of 1–5.3 µm at a resolving power $R \sim 2,700$. Figure 1 shows the geometry of the Io-Jupiter system during our observations. A summary of the observations is provided in Table 1.

The IFU field of view is 3" × 3", with a pixel size of 0.1". With Io's 1.03–1.095" disk, we therefore had a total of 83–94 pixels covering its disk. Io's central meridian longitude was roughly 346°, and the sub-observer latitude 3.3° (see Table 1 for specifics).

JWST uses an up-the-ramp readout mode, where the array is read out non-destructively at intervals defined by the readout pattern. We used the fast readout pattern NRSRAPID (10.737 s/group) for the G235H and G395H grisms, which saves every frame as a group. We used 2 groups per integration. Each integration begins with a detector reset. The G395H data had 2 integrations for each of the 4 dithers; the G235H data had 4 integrations for each of the 4 dithers. For the G140H observations, where Io is fainter, we used the NRSIRS2RAPID (14.589 s/group) readout pattern, which also saves every frame as a group. We used 10 groups per integration, and 4 integrations per exposure, for each of 4 dithers.

<table><caption>**Table 1**
James Webb Space Telescope (JWST) Observations of Io-In-Eclipse, August 2023</caption><thead><tr><th>UT Date/time<br>hr:m:s - hr:m:s</th><th>Disperser/Filter</th><th>Dither</th><th>Exposure<sup>a</sup><br>seconds</th><th>Wavelength<sup>b</sup><br>µm (nm)</th><th>Δ<sup>c</sup><br>AU</th><th>D<sup>d</sup><br>"</th><th>Longitude<sup>e</sup><br>°</th><th>Latitude<sup>e</sup><br>°</th></tr></thead><tbody><tr><td colspan="9"><strong>08/04</strong></td></tr><tr><td>06:39:54–06:53:19</td><td>G395H/F290LP</td><td>1–4</td><td>43</td><td>2.87–5.27 (0.665)</td><td>4.9066</td><td>1.028</td><td>344.7–346.7</td><td>3.26</td></tr><tr><td>06:58:30–07:16:10</td><td>G235H/F170LP</td><td>1–4</td><td>86</td><td>1.66–3.17 (0.396)</td><td>4.9066</td><td>1.028</td><td>347.4–349.9</td><td>3.26</td></tr><tr><td colspan="9"><strong>08/23</strong></td></tr><tr><td>17:44:10–17:54:52</td><td>G140H/F100LP</td><td>1</td><td>584</td><td>0.97–1.89 (0.235)</td><td>4.611</td><td>1.095</td><td>344.0–345.3</td><td>3.36</td></tr><tr><td>17:58:01–18:08:43</td><td>G140H/F100LP</td><td>2</td><td>584</td><td>0.97–1.89 (0.235)</td><td>4.611</td><td>1.095</td><td>346.0–347.5</td><td>3.36</td></tr><tr><td>18:11:38–18:22:20</td><td>G140H/F100LP</td><td>3</td><td>584</td><td>0.97–1.89 (0.235)</td><td>4.611</td><td>1.095</td><td>347.8–349.4</td><td>3.36</td></tr><tr><td>18:25:20–18:36:02</td><td>G140H/F100LP</td><td>4</td><td>584</td><td>0.97–1.89 (0.235)</td><td>4.611</td><td>1.095</td><td>349.8–351.3</td><td>3.36</td></tr></tbody></table>

<footnote><sup>a</sup>Effective exposure time per dither. <sup>b</sup>Wavelength range, with the wavelength spacing (in nm) in brackets. <sup>c</sup>Sub-observer (JWST) distance. <sup>d</sup>Io's diameter. <sup>e</sup>Sub-observer (JWST) W. longitude and latitude.</footnote>
Photons are accumulated continuously during an integration. However, users may delete groups from an integration when starting from the uncalibrated raw data (uncal) to reduce issues with detector saturation. In contrast to the 15 November 2022 observations (de Pater, Lellouch, et al., 2023), this time none of the pixels were saturated since there were, apparently, no energetic eruptions at the time of the August 2023 observations.

The Observatory Scheduling Software used was version 8.7. The raw (uncal) data were calibrated with pipeline version 1.11.3, and CRDS (Calibration References Data System) files CRDS_VER = 11.17.2 and CRDS_CTX = jwst_1,110.pmap for the 4 August and CRDS_CTX = jwst_1,117.pmap for the 23 August data. Two consecutive pipelines are run: Detector Pipeline produces the level 1 data ("rate" and "rateints" files), and Spec2Pipeline produces the level 2 data ("cal" and final images per dither). As for the November 2022 data, we determined the frequency-dependent 1/f noise in each of the three gratings (de Pater, Lellouch, et al., 2023; Trumbo et al., 2023), and subtracted the 1/f noise images from the rate files, before rerunning the standard Level 2 pipeline again to obtain the final calibrated images for each dither. We used the coordinate system "ifualign" in the CubeBuild mode.

The calibration procedure generates calibrated image data cubes in units of MJy/sr. As before (de Pater, Lellouch, et al., 2023), we aligned each set of four dithers using a Fast Fourier transform (FFT) cross-correlation algorithm, and median averaged the data cubes to help remove cosmic ray hits and bad pixels. This resulted in one image data cube in each of the three grisms.

We determined the wavelength-dependent background emission in areas around Io, and subtracted this from the image data cubes; this removed broad emission features due to Jupiter's stray light in the G235H grism at 1.86–1.92 µm and at 2.6–2.8 µm (see Figure 7a in de Pater, Lellouch, et al. (2023)). For the G140H grism, Jupiter's stray light is much worse, particularly outside of its methane absorption bands, as shown in Figure 2a. After (manually) removing cosmic rays from the data cube, we modeled Jupiter's stray light by fitting linear slopes through each row on each image (outside of Io itself) in the spectral data cube, and subtracted this from the original data cube.

To further clean up the image data cubes, we inspected each image and replaced numerous remaining cosmic ray hits and other bad pixels with values at the same pixel location but 1 or 2 spectral channels away in the image data cube. In contrast to the rather smooth variation in wavelength, the pixel-to-pixel variation in intensity across Io is large due to the coarse spatial pixel scale (0.1"), so we could not replace bad pixels with the median of surrounding points, the technique usually used in infrared astronomy (e.g., de Pater, Davies, McGregor, et al., 2014). The resulting spectra are shown in Figures 2b–2d. Thanks to these cleaning techniques, we clearly see the forbidden $a^1\Delta \Rightarrow X^3\Sigma^-$ rovibronic SO emission band at 1.707 µm. Moreover, at shorter wavelengths, at 1.0824 and 1.1308 µm, two narrow lines are visible, which are attributed to forbidden atomic sulfur emissions, [SI], such as seen on the Sun (Swings et al., 1969) and in supernova 1987A (Meikle et al., 1989).

## 2.1. Units

The NIRSpec calibration software produces calibrated spectral radiance cubes in MJy/sr. We converted our image data cubes into photon spectral radiance $I$ (photons s$^{-1}$ cm$^{-2}$ µm$^{-1}$ sr$^{-1}$) (de Pater, Lellouch, et al., 2023):

$$I = L_{\nu} S_{\text{MJy}} \frac{\lambda}{hc} \frac{c}{\lambda^2} = L_{\nu} \times (1.51 \times 10^9 / \lambda), \quad (1)$$

where $S_{\text{MJy}} = 10^{-17}$ erg s$^{-1}$ cm$^{-2}$ Hz$^{-1}$, $hc/\lambda$ is the photon energy, $h = 6.63 \times 10^{-27}$ erg s, and $c/\lambda^2$ converts from power per unit frequency to power per unit wavelength, with wavelength $\lambda$ in µm.

One can also convert the units to flux density per pixel, $F_{\lambda}$, in erg s$^{-1}$ cm$^{-2}$ µm$^{-1}$:

$$F_{\lambda} = L_{\nu} S_{\text{MJy}} \frac{c}{\lambda^2} \times p_{sr} = L_{\nu} \times (3.0 \times 10^{-3} \times p_{sr}/\lambda^2), \quad (2)$$

where $c = 3.00 \times 10^{14}$ µm/s (same length unit as wavelength), and $p_{sr} = 2.35 \times 10^{-13}$, the area of a 0.1" pixel in steradians.

The units can be further converted into the distance-independent radiance, GW µm$^{-1}$ sr$^{-1}$:
<figure><img src="image_4.png" /><figcaption>**Figure 2.** (a) A spectrum of Io's mean radiance in the G140H/100LP filter, before subtracting backgrounds. The spectrum is dominated by Jupiter's stray light. (b) Spectrum of Io's mean radiance in the G140H/100LP filter after subtracting backgrounds, on the same scale as in panel (a). Note that the y-axis in Figure 7 of de Pater, Lellouch, et al. (2023) should be normalized to Io's disk (decreased by a factor of ~100) to get the mean radiance, as shown here. (c) Spectrum of Io's mean radiance as in panel (b), but enlarged by a factor of 10. With the high-resolution setup in G140H, there is a gap in the data at ~1.45 µm. The SO emission band and the [SI] 1.0824 and 1.1308 µm emission lines are indicated. (d) Spectrum of the total intensity of Io after all data cubes were thoroughly cleaned. The different colors of the data refer to the three different grisms (G140H, G235H, and G395H). The light blue spectrum shows the James Webb Space Telescope spectrum taken on 15 November 2022 (from de Pater, Lellouch, et al. (2023)). Note the very different slopes of the spectra, caused by a more vigorous style of resurfacing activity at Kanehekili Fluctus and Loki Patera in November 2022 than in August 2023 (For unit conversions, see Equations 1–3).</figcaption></figure>

$$ \text{Radiance} = F_{\lambda} \times \Delta^{2} \times 10^{-3}/10^{9}, \qquad (3) $$

with Δ being Io's distance from JWST in meters (Table 1). The total radiance from the emitting area on Io's surface in units of GW/µm, assuming Lambertian (diffuse) emission, is obtained by multiplying the value from Equation 3 by a factor of π (Harris, 2013; Wolfe, 1998).

Emission intensity spectra are shown in units of photons s⁻¹ cm⁻² µm⁻¹ sr⁻¹ (Equation 1). We assume that photons from atoms/molecules in Io's atmosphere are emitted isotropically, that is, over 4π steradians. Hence, integrating the spectrum over wavelength (i.e., summed over the spectrum and multiplied by the wavelength spacing), and multiplying by 4π gives the brightness or photon flux in photons s⁻¹ cm⁻²; further multiplication by 10⁻⁹ provides the brightness in kR (kilo-Rayleighs) (Brown et al., 1983).

Note that when spectra are provided in units of photons s⁻¹ cm⁻² µm⁻¹ sr⁻¹, one can average the intensity over a chosen solid angle Ω (e.g., Io's disk) to get the mean radiance over Ω, that is, normalize the brightness to Ω. Alternatively, the method we use is to integrate the flux density (I [Equation 1] × p<sub>sr</sub>) over an area (e.g., Io's disk,
<figure><img src="image_5.png" /><figcaption>Io with JWST/NIRSPEC in August 2023</figcaption></figure>

**Figure 3.** James Webb Space Telescope images of Io in eclipse at 1.4, 1.7, 3.0, and 5.0 µm. The image at 1.7 µm is taken at the center of the SO band (1.705–1.709 µm), and hence contains both thermal volcanic and SO emissions; the latter is minute compared to the thermal emissions. The intensity scale on each panel is chosen such as to show the full range of emissions, using an asinh scaling function. Each circle shows the approximate limb of Io with its North pole oriented upward. The FWHM of the point spread function, at bottom right, varies with wavelength from 0.2 to 0.28", as determined from a stellar point source. The volcanoes Kanehekili Fluctus and Loki Patera are indicated, as well as the potential identification with Gauwa Patera and Sengen Patera. Note that Kanehekili Fluctus is brightest at short wavelengths, but at longer wavelengths Loki Patera is the brightest source. Panels (a) and (b) are from the G140H/F100LP setting on 23 August 2023, and are averaged over 4.7 nm; panel (c) is from G235H/F170LP on 4 August 2023, averaged over 8 nm; panel (d) is from G395H/F290LP on 4 August 2023, averaged over 6.7 nm.

+ a width equal to the extent of the point spread function [PSF] beyond the limb), and divide by Ω (=Io's disk) to normalize the brightness to Io's disk. If the source of emission is unresolved, that is, smaller than the size of the PSF (like a star, or hot spots on Io), one should integrate the flux density over the surface area of the PSF. In such cases we have assumed Ω = p<sub>sr</sub>. In JWST data, the hot spots we have identified are almost certainly sub-pixel. The pixel size is therefore an upper limit for hot spot area. In reality, any gas emitting from the hot spot area will spread further than the boundaries of the hot spot.

**3. Io's Thermal Hot Spot Emissions**

In this section we start with a discussion of the images and spectra of Io's thermal emission, and how we navigate
the image data cubes. We then discuss the volcanoes Kanehekili Fluctus and Loki Patera in detail.

**3.1. Images of Io's Continuum Emission**

Images of Io in eclipse at 1.4, 1.7, 3.0, and 5.0 µm extracted from the image data cubes are shown in
Figures 3a–3d. We determined the location of Io's limb in two ways: Method (i): Using the images in Figure 3, we
<figure><img src="image_6.png" /><figcaption>Keck images in June, August, & December 2023</figcaption></figure>

**Figure 4.** Images of Io taken with NIRC2 on the Keck 2 telescope in 2023, on 22 June, 11 August, and 20 December, at different wavelengths, as indicated. See Table 2 for Io's ephemerides, including the sub-JWST longitudes and latitudes. Io North is up in all frames.

defined a threshold in intensity, and determined the mean position in X,Y pixel coordinates of all spaxels with an intensity above this threshold value (Bockelée-Morvan et al., 2024). This method works well if the full disk is illuminated; however, in Jupiter's shadow Io's emission is mainly localized to thermal sources at the volcanic hot spots. But since there are no zero-values on the disk, the method is expected to work reasonably well. We used values between 0.5% and 10% of the maximum observed intensity, which gave us a typical pixel uncertainty in all maps of $\le 0.02''$. Method (ii): As in our previous paper (de Pater, Lellouch, et al., 2023), we reprojected the images and overplotted them on a USGS map of the satellite (https://astrogeology.usgs.gov/maps/io-voyager-galileo-global-mosaics), while we choose the center of Io such that the location of a particular volcano would agree with that on the USGS map. In our previous paper we anchored the position to that of Loki Patera. We then derived that the second bright spot was Kanehekili Fluctus.

To determine the exact location of the emissions in the JWST data, we turned to observations taken with the near-infrared camera NIRC2 coupled to the Adaptive Optics (AO) system on the Keck 2 telescope (Wizinowich et al., 2000). Figure 4 shows images taken on 22 June, 11 August, and 20 December 2023 (see Appendix A for details on data reduction and analysis). The June data show that Amaterasu Patera was quite active, which might affect the position of Loki Patera's emission in the JWST data since the distance between the two paterae is less than the spatial resolution in the latter data. On 11 August 2023 Kanehekili Fluctus was the brightest source on the disk.

One may debate which of the two methods provides the most accurate center of Io's disk. Kanehekili's latitude as derived from the Keck data with method *i* was $18.2 \pm 0.3^\circ$S, while the longitude varied from $27.5 \pm 0.6^\circ$W at $4.67 \, \mu\text{m}$ to $29.2 \pm 0.9^\circ$W at $3.29 \, \mu\text{m}$, that is suggestive of a slight shift with wavelength. However, no shift in wavelength for this source was detected in the individual JWST image data cubes, that is, no shift was seen across each of the G140H (1.4–1.7 $\mu$m), G235H (1.70–3.0 $\mu$m), and G395H (3.0–5.0 $\mu$m) grisms, though the hot spot positions between the grisms did vary ($25 \pm 3^\circ$W, $17 \pm 1.5^\circ$S in G140H; $25 \pm 3^\circ$W, $13 \pm 2^\circ$S in G235H;
<figure><img src="image_7.png" /><figcaption>**Figure 5.** James Webb Space Telescope maps of Io's thermal emission, after reprojecting the images in Figure 3 and overplotting them on a USGS map of the moon (https://astrogeology.usgs.gov/maps/io-voyager-galileo-global-mosaics). The maps are normalized such that the (maximum intensity)¹/₄ is set equal to 100, which enables one to see both bright and fainter volcanoes. The images in panels (b) and (c) are both centered at 1.707 µm; however, panels (a) and (b) were taken on 23 August with the G140H grism, averaged over 4.7 nm, and panels (c) and (d) on 4 August with the G235H grism, averaged over 8 nm. At the shorter wavelengths the peak emission near Loki Patera is shifted upwards, due most likely to emissions by Amaterasu Patera. Panels (e) and (f) are taken with the G395 grism, averaged over 6.7 nm, where panel (e) is at the same wavelength as panel (d). Note that at the longer wavelengths the emission is peaked at Loki Patera, and at 5 µm it is the strongest emission on the map. The uncertainty in position at disk-center is ≲1° in longitude/latitude (32 km; see text). Several volcanoes are indicated on panel (e).</figcaption></figure>

29 ± 2.5°W, 15 ± 2°S in G395H). Since a real shift should have been noticeable across the individual grisms, we decided to use method (ii) with Kanehekili's position fixed at 28.5°W and 18.2°S to determine the location of Io's center and limb in the JWST data. Using method (ii) on the Keck image from 20 December 2023, with the above position for Kanehekili, gives a position for Loki Patera at 306.5°W and 13.6°N, that is, on the Eastern side of the horseshoe-shaped lake. The final reprojected JWST images, superposed on the USGS map, are displayed in
<table><caption>**Table 2**

Keck Observations of Io in 2023</caption><thead><tr><th>UT time<br>year month/day hr:m</th><th>Filter<sup>a</sup></th><th>Wavelength</th><th>Δ<sup>b</sup><br>AU</th><th>CML<sup>c</sup><br>°</th><th>Latitude<sup>c</sup><br>°</th><th>Volcano</th></tr></thead><tbody><tr><td>2023 06/22 15:18</td><td>L'</td><td>3.78</td><td>5.50</td><td>313.7</td><td>3.07</td><td>Loki P.</td></tr><tr><td>2023 06/22 15:09</td><td>Ms</td><td>4.67</td><td>5.50</td><td>312.4</td><td>3.07</td><td>Loki P.</td></tr><tr><td>2023 06/22 15:12</td><td>PAH</td><td>3.29</td><td>5.50</td><td>312.9</td><td>3.07</td><td>Loki P.</td></tr><tr><td>2023 08/11 15:34</td><td>L'</td><td>3.78</td><td>4.79</td><td>44.3</td><td>3.32</td><td>Kanehekili F.</td></tr><tr><td>2023 08/11 15:38</td><td>Ms</td><td>4.67</td><td>4.79</td><td>44.9</td><td>3.32</td><td>Kanehekili F.</td></tr><tr><td>2023 08/11 15:41</td><td>PAH</td><td>3.29</td><td>4.79</td><td>45.3</td><td>3.32</td><td>Kanehekili F.</td></tr><tr><td>2023 12/20 05:23</td><td>L'</td><td>3.78</td><td>4.32</td><td>344</td><td>3.14</td><td>Kanehekili F., Loki P.</td></tr><tr><td>2023 12/20 09:30</td><td>L'</td><td>3.78</td><td>4.32</td><td>19.4</td><td>3.14</td><td>Kanehekili F.</td></tr><tr><td>2023 12/20 09:38</td><td>Ms</td><td>4.67</td><td>4.32</td><td>20.8</td><td>3.14</td><td>Kanehekili F.</td></tr></tbody></table>

Figure 5. Based on the Keck data, the overall uncertainty in position at disk-center is estimated to be ≲1° (32 km) in latitude/longitude on Io.

As shown in Figures 3 and 5, Kanehekili Fluctus is the brightest source on the maps at wavelengths 1–3 µm, while Loki Patera is brightest at 5 µm. At shorter wavelengths the peak intensity at Loki Patera is shifted northwards, due most likely to emission from Amaterasu Patera (Figure 4).

## 3.2. Io's Thermal Emission Spectra

Spectra of the radiance of Loki Patera and Kanehekili Fluctus are shown in Figure 6. Since the thermal emission from these volcanoes is unresolved in our images, we determined their brightness by integrating the flux density in a circular aperture with a radius of 1 pixel (0.1"), using the apphot.pro program in IDL (Perrin, 2013), after increasing the size of each image frame in the data cube by a factor of 5 to create sub-spaxels at a 0.02"-sized grid. By using aperture-summed spectra instead of single-pixel spectra we minimized the oscillations in individual spaxel spectra due to the undersampling of the PSF. Since we integrate the intensity of the volcanoes over an aperture much smaller than the size of the PSF, the numbers were corrected for the loss of flux due to the extent of the PSF (de Pater, Davies, McGregor, et al., 2014; Gibbard et al., 2005). We determined this loss by measuring the intensity of the star P330-E (Bohlin & Landolt, 2015) taken with the same grisms and the same small circular aperture. The ratio of this intensity to the total intensity of the star was ~40% at the shortest wavelengths, down to ~25%–30% at the longest wavelengths. The measured intensities of Loki Patera and Kanehekili Fluctus were corrected for this wavelength-dependent loss of flux. These data were further corrected for the emission angle effect, that is, they were divided by cos(θ), with θ the emission angle, to correct for the geometric foreshortening effect when a flat surface on a sphere is not at disk-center. This resulted in multiplication factors of ~1.45 for Kanehekili Fluctus, and ~1.3 for Loki Patera. All spectra were further normalized to the 2023 August 23 date.

One may notice a slight jump between the G140H/F100LP and G235H/F170LP grisms. We have not attempted to shift spectra such as to make them agree in the overlapping region, since it is not clear which grism has the best calibration. We also note that the background subtraction in the G140H grism (Section 2 and Figure 2a) may have affected the spectra somewhat, and the two spectral regions are not sampled concurrently in time while volcanic eruptions are episodic.

As shown in Figure 6, the shape of the 2022 and 2023 spectra are very different: Kanehekili Fluctus was extremely active in November 2022, and Loki Patera had just entered a new brightening phase (de Pater, Lellouch, et al., 2023). The high levels of activity are supported by NASA Infrared Telescope Facility (IRTF) data (Shah et al., 2023) and observations by the *Juno* Jovian Infrared Auroral Mapper (JIRAM) (Davies et al., 2024; Perry et al., 2025). In the following subsections, we compare the 2023 spectra with those obtained in 2022. In the first section, we use the same method as used in de Pater, Lellouch, et al. (2023), that is, we determine good matches to the spectra by eye, using 4–5 distinct blackbody curves and without constraining total area. In the second section we

<footnote><sup>a</sup>L' filter spans the wavelengths range 3.426–4.126 µm; Ms filter spans 4.549–4.790 µm; PAH filter spans 3.263–3.318 µm. <sup>b</sup>Geocentric distance. <sup>c</sup>Sub-observer W. longitude (CML, central meridian longitude) and latitude.</footnote>
<figure><img src="image_8.png" /><figcaption>**Figure 6.** Spectra of the flux density of Kanehekili Fluctus panels (a) and (b) and Loki Patera panels (c) and (d). The August 2023 data are shown in black; the light blue curves panels (a) and (c) show the James Webb Space Telescope spectra from November 2022, scaled appropriately to the 23 August 2023 date (de Pater, Lellouch, et al., 2023). The spectra for Kanehekili Fluctus and Loki Patera were smoothed over wavelength, using a boxcar function of 5 pixels. These spectra were taken at the peak intensity on the maps. As discussed in the text, both volcanoes are unresolved and the spectra represent the total flux density of their thermal emissions. Superposed are blackbody spectra at different temperatures, together with the relevant areas (indicated by a radius, assuming a circular area) as indicated. The red line shows the summation of these curves. On the left panels (a) and (c) the entire observed spectra are shown (0.97–5.53 µm); on the right panels (b) and (d) we show an enlarged view of the 0.97–2.5 µm region. We note that the spectra do not exactly line up in the region of overlap, which coincides with the SO emission band. We did not try to shift one of the spectra to align them. Judging from the blackbody fits, the “alignment” appears to be quite good.</figcaption></figure>

determine multi-Temperature, multi-Area fits (with multi = three) to the 2022 and 2023 spectra of Loki Patera where we constrain the total area of the model components to the known emitting area of Loki Patera, which exhibits a strong control on model fits. In the third subsection we focus on the eruption at Kanehekili, and apply the Io Flow Model (IFM (Davies, 1996)) to the data. In the last subsection we discuss the location of the eruptions in the Kanehekili area, and construct a timeline of observations with *Juno/JIRAM* and *Keck* data.

### 3.3. Fits to Thermal Emission Spectra Using Multiple Blackbody Curves

Superposed on the observed spectra in Figure 6 we show several blackbody curves, the sum of which matches the observations quite well. Although the temperature of a cooling lava flow should be a continuum from the eruption temperature down to the background equilibrium value, one can combine a finite number of blackbody curves between 100 and ~1,500 K (assuming a basalt composition), each having a unique spectral shape, to find a good match to the JWST spectra. As shown in de Pater, Lellouch, et al. (2023), this simplified discretization of temperatures provides acceptable fits, with typical uncertainties on the derived areas of order 10%. We initially use the same temperature discretization values as in 2022 for the sake of comparing like with like, and indicate these with the derived areas in Figure 6, as well as in Table 3. As shown, the spectra are consistent with cooling lava flows, with the hottest lavas covering the smallest areas, rapidly increasing in size at lower temperatures. We note that the highest temperatures are obtained by matching the observed spectral shape at the shortest wavelengths, while the observed spectra at longer wavelengths are sensitive to the lower temperatures.
<table><caption>**Table 3**
Comparison in Temperature and Area From Blackbody Fits in November 2022 and August 2023</caption><thead><tr><th>Temperature K</th><th>Area (November 2022) km<sup>2</sup></th><th>Area (August 2023) km<sup>2</sup></th></tr></thead><tbody><tr><td colspan="3"><strong>Loki Patera</strong></td></tr><tr><td>1,200</td><td>0.06</td><td>—</td></tr><tr><td>900</td><td>2.0</td><td>0.50</td></tr><tr><td>600</td><td>32</td><td>8</td></tr><tr><td>400</td><td>774</td><td>452</td></tr><tr><td>300</td><td>5,026</td><td>—</td></tr><tr><td>280</td><td>—</td><td>19,100</td></tr><tr><td colspan="3"><strong>Kanehekili Fluctus</strong></td></tr><tr><td>1,200</td><td>0.26</td><td>0.17</td></tr><tr><td>900</td><td>3.1</td><td>1.5</td></tr><tr><td>600</td><td>69</td><td>28</td></tr><tr><td>400</td><td>876</td><td>314</td></tr><tr><td>300</td><td>2,027</td><td>5,026</td></tr></tbody></table>

In November 2022, the highest temperature component modeled at Kanehekili Fluctus was 1,200 K, covering an area of 0.26 km². Nine months later, when activity at Kanehekili Fluctus had subsided considerably, a similar temperature component was still detectable, but the area had decreased by almost 40% to 0.17 km². The area decreased more for the lower temperature components (a factor of 2.8 at 400 K, from 876 to 314 km²). At still lower temperatures, though, the model component area had increased by a factor of ~2.5, as manifested by the 300 K blackbody curve fit (from ~2,000 to 5,000 km²; Table 3). This trend at low temperatures is consistent with the spreading and cooling of lava flow surfaces between observations.

Loki Patera had just entered a brightening phase when JWST observed it in November 2022 (de Pater, Lellouch, et al., 2023), caused presumably by a new overturn in its lava lake (de Kleer & de Pater, 2017; Matson et al., 2006; Rathbun & Spencer, 2006). Note the very different spectral shapes for Kanehekili Fluctus and Loki Patera. For Kanehekili Fluctus, the difference between the 2022 and 2023 spectra continues to increase with increasing wavelength (up to 5 µm), while for Loki Patera the difference is largest near 3.5 µm, and becomes almost negligible near 5 µm. The highest temperature we could match to the short-wavelength data for Loki Patera in 2023, even though our new spectral range goes all the way down to 1 µm, was 900 K. Both the 600 and 900 K areas in 2023 are a factor of 4 smaller than in 2022; at

400 K the area is roughly a factor of 2 smaller in 2023. In contrast to the fits at shorter wavelengths, at wavelengths beyond 4 µm the 2023 spectrum can be matched in several ways by choosing temperatures between 200 and 300 K in addition to the 400–900 K temperatures. For example, a 300 K temperature over ~5,000 km² together with either a 200 K over ~125,500 km² or 250 K over ~30,000 km² (or values in between) match the observations well. An equally good fit is found by replacing the combination of these two temperatures by a single temperature between 270 K (area 28,000 km²) and 280 K (area 19,100 km²). Given the finite extent of the patera (Section 3.4), the 280 K temperature provides the best choice.

Assuming an ongoing eruption with a constant areal coverage rate (which can be either an expanding area of lava flows, or the replacement of the crust on a lava lake), a surface temperature distribution forms where the rate of cooling drops with time, and the evolution of the cooling lava surfaces at a volcano can be calculated, as we do in the next two subsections.

### 3.4. Loki Patera: Three-Temperature, Three-Area Constrained Fits

The emitting area of Loki Patera is strongly constrained by topography. The exposure of new lava and the replacement of old lava is confined to the floor of the patera, excluding the higher-albedo “island” and some small areas close to the “island.” These higher-albedo regions have hardly changed over the decades since Loki Patera was first imaged by *Voyager 1* in 1979 (Smith et al., 1979), as seen from images presenting the same appearance in *Galileo* (Turtle et al., 2004), *New Horizons* (Spencer et al., 2007) and *Juno* data (Ravine et al., 2024). In the infrared, however, Loki Patera is highly variable, with total thermal emission generally varying between 6 and 16 TW (Davies et al., 2012; de Kleer & de Pater, 2017; de Pater, de Kleer, Davies, & Ádámkovics, 2017), with an average of ~9 TW (Matson et al., 2006). The analysis of Loki Patera's thermal emission data from multiple instruments strongly suggests the patera is filled with a lava lake that periodically overturns (Davies, 2003; de Kleer & de Pater, 2017; de Kleer et al., 2017; Matson et al., 2006; Rathbun et al., 2002); the most recent, though indirect, evidence in support of Loki Patera being a lava lake is shown in Figure 6, discussed in the previous subsection. *Galileo* and *Juno* images showed that Loki Patera is covered by a relative smooth glassy surface that exhibits specular reflectance, as do some other paterae on Io (Radebaugh et al., 2024; Turtle et al., 2004). Resurfacing models predict and spacecraft observations reveal that all areas of the emitting surface—the low albedo area of the patera floor—are always at temperatures near or above ~250 K (Davies, 1996; Matson et al., 2006; Rathbun & Spencer, 2006), yielding spectral radiance detectable in the wavelength range covered by infrared-sensitive instruments (Davies et al., 2010).
<figure><img src="image_9.png" /><figcaption>**Figure 7.** (a) Three-temperature, three-area fits to Loki Patera spectra from November 2022 and August 2023. The total emitting area was fixed at 21,500 km², the area of the low-albedo unit on the patera floor (Veeder et al., 2011). Temperatures, areas, and total thermal emission are shown in Table 4. (b) Comparison of the James Webb Space Telescope spectra in panel (a) with four *Galileo/NIMS* spectra from Davies and Veeder (2023). These spectra show how thermal emission from Loki Patera evolves and spectral shape changes during a brightening event (Matson et al., 2006). G7 (April 1997—curve 5) was obtained during the height of a brightening event; and C9 (June 1997—curve 6) was toward the end of the resurfacing event. Thermal emission then waned, as shown by E16 (July 1998—curve 7), finally reaching a quiescent, low thermal emission period dominated by passive cooling of replaced lava lake surface by C22 (August 1999—curve 8). NIMS data are shown as stars, and the solid lines are two-area, two-temperature model fits to the NIMS data (see Davies and Veeder (2023)). The model fit for E16 has been refined from that published in Matson et al. (2006).</figcaption></figure>

The low-albedo area is $2.15 \times 10^4$ km² (Veeder et al., 2011). This provides a further constraint on fitting the JWST spectra (as with *Galileo* NIMS data, Davies, 2022; Davies et al., 2012), as the combination of different areas and temperatures must not only fit the JWST spectra, but the areas of the different model components must combine to match the total emitting area. Figure 7 shows a three-temperature, three-area fit to both the November 2022 and August 2023 JWST data, where the total emitting area is constrained to be equal to 21,500 km². Table 4 shows the temperature and areas for the best (least-squares) model fits. We stress that, like in the previous subsection, the temperatures are expressions of a complex temperature distribution across the floor of the patera, and are the effective temperatures for the different regions. Resurfacing models yielding the evolving thermal emission (de Kleer & de Pater, 2017; Matson et al., 2006) predict a broad range of surface temperatures and ages.
<table><caption>**Table 4**
Three-Temperature, Three-Area Fits to Loki Patera</caption><thead><tr><th>Component</th><th>Area km<sup>2</sup></th><th>Temperature K</th><th>Thermal emission TW</th><th>Time to cool<sup>a</sup> days</th></tr></thead><tbody><tr><td colspan="5">November 2022</td></tr><tr><td>Component 1</td><td>21,301</td><td>272</td><td>6.61</td><td>266</td></tr><tr><td>Component 2</td><td>197</td><td>511</td><td>0.76</td><td>117</td></tr><tr><td>Component 3</td><td>1.95</td><td>910</td><td>0.08</td><td>7 min</td></tr><tr><td>Totals</td><td>21,500</td><td></td><td>7.45</td><td></td></tr><tr><td colspan="5">August 2023</td></tr><tr><td>Component 1</td><td>21,038</td><td>280</td><td>7.33</td><td></td></tr><tr><td>Component 2</td><td>460</td><td>393</td><td>0.62</td><td></td></tr><tr><td>Component 3</td><td>2.3</td><td>760</td><td>0.04</td><td></td></tr><tr><td>Totals</td><td>21,500</td><td></td><td>8.00</td><td></td></tr></tbody></table>

A three-temperature, three-area model broadly reproduces this more complex temperature-area distribution, though may “miss” the highest temperature components contributing to thermal emission at the shortest wavelengths (Tables 3, and 4 in this paper, and Figure 6 in de Pater, de Kleer, Davies, and Ádámkovics (2017)).

From the perspective of modeling the heat loss from a lava lake, which is unlike a lava flow in as much as a lava flow has a finite thickness with heat being lost from upper and lower surfaces, a lava lake can be treated as having only one (upper) surface through which heat is lost. A reasonable assumption is that magma beneath the forming surface crust is isothermal and, for the purposes of modeling (Davies et al., 2005; Matson et al., 2006), of infinite depth—a “semi-infinite space.” Table 4 shows that the bulk of the patera floor in 2022 is at an effective temperature (Veeder et al., 2009) $T_{eff}$ = 272 K, suggesting an age (time since exposure) of 266 days. This agrees well with the cooling timing since the end of the last resurfacing event (Davies et al., 2005; de Kleer & de Pater, 2017; Matson et al., 2006). Excess thermal emission at short wavelengths, when compared with the later 2023 data, is most likely the result of part of the patera floor being in the process of replacement. These new, hotter areas shift the peak in thermal emission to shorter wavelengths, represented here by $T_{eff}$ ~ 510 K (with an age of 117 days), combined with a very small area, ~2 km², at ~910 K. This latter area represents the youngest lava, exposed during the resurfacing process at locations where the crust breaks up at lake margins, and through cracks in the cooler lava surface crust. As shown in Section 3.3, some cracks expose even hotter temperatures (1,200 K over 0.06 km²).

Table 4 shows the thermal emission from each model component calculated using the Stefan-Boltzmann Law given in Equation 4 (e.g., Veeder et al., 2009). Table 4 also shows total thermal emission.

$$P = A\epsilon\sigma T^4 \quad (4)$$

In Equation 4, *T* is the component temperature, *A* the area, *ε* the emissivity (taken to be unity), and *σ* is the Stefan-Boltzmann constant. The total emitted power from Loki Patera was 7.45 TW in 2022 and 8.0 TW in 2023. Even though the 2022 data were taken at the beginning of a brightening, Loki Patera actually emits less power than during the quiescent time in 2023, when the power output is dominated by emission from large, relatively cool areas. In 2022, the bulk of the surface of Loki Patera was at cooler temperatures than in 2023 when the lava crust had not yet reached its lowest temperatures before the next brightening event started. When comparing the power at only high temperatures, at or above 600 K, it is four times higher in 2022 (0.32 TW) than in 2023 (0.08 TW), as derived from the multi-blackbody fits in the previous subsection. We note that the total power estimates for 2022 and 2023 are still a factor of 2 below the maximum power at the height of a brightening event. Figure 7b shows a comparison of the JWST spectra with NIMS spectra taken during different times at brightening events. JWST would not have been able to capture Loki Patera's thermal spectrum at the height of a brightening event due to saturation effects.

<footnote>ªTime for newly-exposed lava to cool to this temperature, assuming Loki Patera contains a lava lake.</footnote>
<figure><img src="image_10.png" /><figcaption>**Figure 8.** Spectra of the Kanehekili Fluctus region with superposed model outputs from the Io Flow Model (IFM) (Davies, 1996, 2007). Panel (a) shows the 15 November 2022 data (black line), and panel (b) shows the combined 4 and 23 August 2023 data (black line). The red curves (1) are for IFM component 1; the green curves (2) for IFM component 2; and the blues curves (3) are the sum of the two components, showing the best fit to the James Webb Space Telescope spectra. Table 5 provides a summary of the fitted parameters.</figcaption></figure>

### 3.5. Kanehekili Fluctus: Spectral Fits Using the Io Flow Model (IFM)

In addition to the multi-temperature fit discussed above (Section 3.3), we model the Kanehekili spectra with the IFM, which determines the thermal emission spectrum from the surface of a steadily spreading, and cooling, basaltic lava flow (Davies, 1996, 2007). In this model, we adopt an eruption temperature of 1,475 K, the liquidus temperature of basalt, and calculate the solidification and cooling behavior of the lava based upon the thermal and physical characteristics of basalt over time. We then calculate the distribution of resulting temperatures on the surface, and determine the infrared thermal emission spectrum from this model as a function of time. We vary the duration of the event and the area covered by new lava in an iterative way to fit the observations. To get better fits to the data we use two model components rather than one, where the eruption temperature of each component starts at 1,475 K. Table 5 provides a summary of the fitted parameters, and Figure 8 shows the results.

As shown, the Kanehekili Fluctus region was extremely active on 15 November 2022. The area of flows thermally detectable by JWST covered about 990 km². The model fit is consistent with the bulk of the thermal emission emanating from extensive flows with an insulating upper crust, and a temperature distribution down to 375 K, suggesting the flows were emplaced over a period of ~12 days. In addition to this flow unit, there is a much smaller but much hotter component (component 2) which is interpreted to represent the sum of the active vent, breakouts, and cracks in the crust through which extremely hot material is exposed. This component has a temperature range from the eruption temperature down to 875 K, the lower temperature reached in just about 7 min. The total area of this component is 2.1 km². If this were a lava flow in its own right, it would suggest a rapid rate of areal coverage at an outflow velocity of 5,000 m² s⁻¹. Practically, this model component makes up only ~0.2% of the total flow area. Note that the fits with multiple blackbody curves revealed small areas at even higher temperatures (0.06 km² at ~1,200 K), while the areal extent of components at other temperatures agrees quite well between the models, validating the application of the IFM model. The 2-to-5 (2:5) µm spectral radiance ratio is 0.06, consistent with lava flow fields with well-developed insulating surface crusts elsewhere on Io (Davies et al., 2010; de Pater, Davies, Ádámkovics, & Ciardi, 2014).

By August 2023, there was still a small, highly active area present, but it was roughly half the size of that in November 2022 (1.22 km²). This was evidence that the eruption was still ongoing. The total area had increased roughly fourfold to over 4,300 km². So the “hot” component comprised only 0.03% of

the total area, almost an order of magnitude less than 9 months earlier. The oldest areas of crust had dropped in temperature to 275 K, implying an age of ~160 days, a little over five months (Davies et al., 2010), to cool from the initial eruption temperature.

The changing shape of the thermal emission spectrum over time is manifested in a drop in the 2:5 µm spectral radiance ratio from 0.15 to 0.12 and a decrease in the magnitude of thermal emission at JWST wavelengths, while the total modeled thermal emission increased from 2.2 TW in November 2022 to 2.7 TW in August 2023. This is consistent with the emplacement of flows dominated by insulating surface crusts, and an evolving temperature-area distribution that over time moves the peak of thermal emission to longer wavelengths, increasingly beyond 5 microns.
<table><caption>**Table 5**

Fits to the Kanehekili Fluctus Region Using the Io Flow Model</caption><thead><tr><th>Component</th><th>Area<sup>a</sup><br>km<sup>2</sup></th><th>Temperature<sup>a</sup><br>K</th><th>Total power<br>TW</th><th>Time to cool<sup>b</sup><br>hrs</th><th>Flow rate<br>m<sup>2</sup>/s</th></tr></thead><tbody><tr><td colspan="6"><strong>November 2022</strong></td></tr><tr><td>Component 1</td><td>990</td><td>375</td><td>2.0</td><td>280</td><td>985</td></tr><tr><td>Component 2</td><td>2.1</td><td>875</td><td>0.1</td><td>0.12</td><td>5,000</td></tr><tr><td>Totals</td><td>992</td><td></td><td>2.1</td><td></td><td></td></tr><tr><td colspan="6"><strong>August 2023</strong></td></tr><tr><td>Component 1</td><td>4,362</td><td>275</td><td>2.6</td><td>3,900</td><td>312</td></tr><tr><td>Component 2</td><td>1.3</td><td>875</td><td>0.06</td><td>0.12</td><td>3,000</td></tr><tr><td>Totals</td><td>4,363</td><td></td><td>2.7</td><td></td><td></td></tr></tbody></table>

### 3.6. Kanehekili Region: Location of Emissions and Temporal Evolution

Thermal emission in the Kanehekili region of Io was, prior to Juno, generally attributed to volcanic activity at Kanehekili Fluctus. Thermal emission was detected from Kanehekili Fluctus by Juno's JIRAM at 33.1°W, -18.1°N as early as May 2017 (Davies et al., 2024; Perry et al., 2025). Thermal emission from an eruption site to the east of Kanehekili Fluctus centered on 24°W, -18.1°N was first detected in JIRAM data in early 2019 (Davies et al., 2024; Perry et al., 2025) and by Keck in 2020 and 2021 (discussed below). However, it was not until 3 February 2024 that Juno's JunoCam imaged the location of the hot spot, now known as Guaraci Fluctus, which is shown in Figures 9a and 9b. This feature was not observed in eruption by Galileo (1996-2002). As noted by Perry et al. (2025), some small, pre-Voyager lava flows are identified here on the USGS Io geological map (D. A. Williams et al., 2011) used in Figures 5, 15, and 20. The last time the Kanehekili region was imaged at visible wavelengths prior to the Juno mission was 28 February 2007 during the New Horizons flyby (image lor_0034940539_0 × 630_sci_1). New Horizons did not observe any thermal emission at Guaraci Fluctus in the infrared.

We superpose the thermal emission JWST data from 15 November 2022 (de Pater, Lellouch, et al., 2023) and 23 August 2023 on these images (Figure 9) in panels (c) and (d) (panels (e) and (f) show the SO emissions at the center of the emission band, discussed in Section 4.1). The emissions both in 2022 and 2023 seem to originate from a region in between Kanehekili Fluctus and Guaraci Fluctus.

To further investigate the location, we turned to near-IR observations taken with NIRC2 on the Keck 2 telescope. We used observations that were published before (Cantrall et al., 2018; de Kleer, de Pater, Molter, et al., 2019), and Keck data taken since 1 July 2018. The data are summarized in Table 6 (see Appendix A for details). Figure 10 shows the location of emissions superposed on a map of the Kanehekili region. In 2020-2021 the emissions, although quite faint, clearly arose from Guaraci Fluctus (see Figure 11), whereas in 2022-2023 it originated closer to Kanehekili Fluctus (Figure 4). We note that, although the positional accuracy is good to within ~0.5-1° on Io, the PSF (full width at half maximum) is 4-5°. Hence, we do not resolve the Kanehekili region, but the center location of the emission may move around, depending on the exact source(s) and strength of the emissions.

As detailed by Perry et al. (2025), the Kanehekili region has been observed with Juno/JIRAM numerous times since 2017. Up to orbit PJ55 in October 2023, the spatial resolution with JIRAM ranged from 151 km/pixel (PJ35, 21 July 2021) to 13 km/pixel (PJ55, 15 October 2023). Spectral radiance from Kanehekili Fluctus and Guaraci Fluctus could easily be separated (Perry et al., 2025). In Figure 12 we show a time line of the spectral radiances of the Kanehekili and Guaraci Fluctus regions as observed with Juno/JIRAM (from Perry et al. (2025)), Keck and JWST. While before 2022 Guaraci Fluctus appeared to be brighter than Kanehekili Fluctus, as seen by both JIRAM and Keck, this changed after 2022, when most of the thermal emission originated (again) closer to Kanehekili Fluctus (Figure 10).

When compared with the Keck and JIRAM data (Figure 12), the multiwavelength spectral radiance from the Kanehekili region measured by JWST in November 2022 was at its maximum value as observed between 2016 and March 2024. Prior to 2022, activity at both Kanehekili locations was low, with Guaraci Fluctus being the brighter of

<footnote><sup>a</sup>Areal size and temperature of the oldest surface area. <sup>b</sup>Age of the oldest surface area, that is, time to cool to the given temperature.</footnote>
<figure><img src="image_11.png" /><figcaption>**Figure 9.** (a) JunoCam image 58C00025 (JNCE_2024034_58C00025_V01), taken 3 February 2024 at 17:50 UT from an altitude of 2,481.6 km (jovicentric distance of 4,310.6 km). The sub-spacecraft point is 10.797°S, 0.575°W. The phase angle is 121.5°. The pixel scale is 1.67 km (NASA/SwRI/MSSS/Jason Perry; Ravine et al., 2024; Perry et al., 2025). (b) Cylindrical projection of the JunoCam image in panel (a). Kanehekili Fluctus and the new eruption site, Guaraci Fluctus, are indicated (NASA/SwRI/MSSS/Jason Perry; Ravine et al., 2024; Perry et al., 2025). (c) Io's thermal emission at 1.7 µm from 15 November 2022 superposed on the projected JunoCam map from panel (b). (d) Io's thermal emission at 1.7 µm from 23 August 2023 superposed on the JunoCam map. (e) Map of the 15 November 2022 SO emission, integrated over 1.705–1.709 µm, superposed on the JunoCam map. (f) Map of the 23 August 2023 SO emission, integrated over 1.705–1.709 µm, superposed on the JunoCam map.</figcaption></figure>

the two at ~8–10 GW µm⁻¹ from 2019 through 2021 at 4.8 µm, and almost disappearing in 2022. Keck observations showed an up-tick in emission from Kanehekili Fluctus in mid-2022, up to ~40 GW µm⁻¹, culminating in the JWST observation in November 2022 of 175 GW µm⁻¹. As charted by Keck and JIRAM this eruption was followed
<table><caption>**Table 6**

Keck Observations of Kanehekili in 2020–2024<sup>a</sup></caption><thead><tr><th>UT time<br>yr month day</th><th>CML<sup>b</sup><br>°</th><th>W. Long.<br>°</th><th>Lat.<br>°</th><th>Intensity 3.78 µm<br>GW sr<sup>-1</sup> µm<sup>-1</sup></th><th>Intensity 4.67 µm<br>GW sr<sup>-1</sup> µm<sup>-1</sup></th></tr></thead><tbody><tr><td>2020 09 08</td><td>351</td><td>24.6 ± 0.5</td><td>-18.3 ± 1.0</td><td>1.4 ± 0.2</td><td>2.25 ± 0.2</td></tr><tr><td>2021 04 17</td><td>338</td><td>21.6 ± 1.5</td><td>-16.2 ± 1.5</td><td>-</td><td>1.2 ± 0.2</td></tr><tr><td>2021 07 20</td><td>53</td><td>23.0 ± 1.0</td><td>-17.2 ± 1.0</td><td>1.34 ± 0.4</td><td>2.6 ± 0.3</td></tr><tr><td>2022 08 15</td><td>54</td><td>28.4 ± 0.4</td><td>-17.2 ± 0.3</td><td>9.5 ± 1.0</td><td>12.5 ± 1.3</td></tr><tr><td>2022 09 07</td><td>33</td><td>30.2 ± 1.0</td><td>-17.5 ± 0.3</td><td>11 ± 3</td><td>9 ± 3</td></tr><tr><td>2022 09 12</td><td>342</td><td>28.6 ± 0.9</td><td>-17.7 ± 0.5</td><td>10 ± 1</td><td>14.5 ± 1.5</td></tr><tr><td>2023 08 11</td><td>45</td><td>28.5 ± 0.9</td><td>-18.2 ± 0.3</td><td>10.7 ± 1.5</td><td>16 ± 2</td></tr><tr><td>2023 12 20</td><td>20</td><td>31.2 ± 1.0</td><td>-17.0 ± 1.0</td><td>7 ± 1</td><td>11 ± 2</td></tr><tr><td>2024 03 02</td><td>341</td><td>33.7 ± 1.0</td><td>-18.0 ± 0.3</td><td>3.0 ± 0.7</td><td>5.8 ± 0.6</td></tr></tbody></table>

by a gradual decrease over the following months, up to March 2024. We note that the *Juno/JIRAM* data in 2023 should be treated with caution, since they were taken at large emission angles where the emission may be enhanced or partially blocked by topography (de Kleer & de Pater, 2017; Mura et al., 2024). In addition, the 2023 JWST data may have been affected by emissions from nearby Janus and Uta Paterae, and hence might be regarded as an upper limit. In addition to the thermal data from these three instruments, JunoCam data provide visible context (Figure 9), showing lava flow emplacement at both Kanehekili Fluctus and Guaraci Fluctus. At Kanehekili Fluctus this mode of eruption is consistent with episodic eruptions emplacing large lava flows. This scenario was previously proposed from analyses of Kanehekili Fluctus data in 2010. Activity in the region at that time included an extremely powerful eruption where the 5-µm spectral radiance peaked at nearly 500 GW/µm (de Pater, Davies, McGregor, et al., 2014).

We note that the analysis above could only be accomplished by combining data from four instruments, which yield a more complete picture of activity in the Kanehekili region than any instruments alone could have provided us.

## 4. Images and Spectra of the Forbidden SO Emission Band

While the previous section focused on Io's hot spot thermal emissions from its surface, we now switch to gaseous SO and S emissions from the satellite's atmosphere (Sections 4 and 5). To separate these emissions from Io's thermal hot spots, we follow the method outlined by de Pater, Lellouch, et al. (2023). For the G235H grism we modeled the thermal continuum emission by fitting a second-order polynomial through each spaxel in the image data cube. For the G140H grism we used third order polynomials in segments containing the line profiles: 1.02–1.11, 1.11–1.18, and 1.63–1.79 µm. After subtracting these polynomial fits from each spaxel in the image data cube, we further improved the SO emission complex by subtracting an image of the remaining thermal emissions. We constructed this image by median averaging frames over 1.661–1.684 and 1.745–1.769 µm. We then subtracted the average of these two frames from the SO segment (1.63–1.79 µm) of the data cube.

The resulting disk-averaged SO and S emission spectra are displayed in Figure 13. We show the SO spectrum both for the G235H (pink) and G140H (black) data. The spectra agree very well, but clearly the G140H spectrum has a much higher signal-to-noise ratio (SNR) than the G235H spectrum, due to the ~7 times longer integration time and higher detector efficiency (https://jwst-docs.stsci.edu/jwst-near-infrared-spectrograph/nirspec-performance#gs_c.tab=0). We therefore limit our SO discussions in this paper to the G140H data.

<figure><img src="image_12.png" /><figcaption>**Figure 10.** Locations of thermal emission sources at Kanehekili Fluctus (left of image center) and Guaraci Fluctus (centered at a longitude of ~24°W) as observed with the Keck Telescope since 2004, plotted on a JunoCam observation 58C00025 obtained in 2024 (Perry et al., 2025). The data between 2004 and 2017 were taken from Cantrall et al. (2018) and de Kleer, de Pater, Molter, et al. (2019); only data with error bars less than 3° were plotted. The observations taken between 2020 and 2023 are from this paper.</figcaption></figure>

<footnote><sup>a</sup>Kanehekili was not detected on 6 March 2020; its last detection was on 27 May 2017 (de Kleer, de Pater, Molter, et al., 2019). <sup>b</sup>Central meridian W. longitude at the time of the observations.</footnote>
<figure><img src="image_13.png" /><figcaption>**Figure 11.** Images of Io taken with NIRC2 on the Keck 2 telescope in 2020 and 2021 (see Table 6). Io North is up in all frames. Note that, in contrast to Figure 4, Kanehekili does not line up with Janus Patera and Uta Patera; in these 2020–2021 images we detect Guaraci Fluctus, as discussed in the text.</figcaption></figure>

<figure><img src="image_14.png" /><figcaption>**Figure 12.** Time line of spectral radiance from the Kanehekili region as measured by Juno/JIRAM (Perry et al., 2025), the Keck telescope, and the James Webb Space Telescope. Panel (a) shows the spectral radiance at 3.5 & 3.78 µm; panel (b) shows the spectral radiance at 4.67 & 4.8 µm. Juno/JIRAM resolved the hot spots, so spectral radiance from Kanehekili Fluctus and Guaraci Fluctus are shown separately. Squares and circles (Juno/JIRAM) surrounded by black lines indicate data at emission angles >75°, and should be treated with caution.</figcaption></figure>
<figure><img src="image_15.png" /><figcaption>**Figure 13.** (a) Spectrum of the forbidden 1.707 µm SO emissions averaged over Io's disk. The pink line shows the spectrum in the G235H/F170LP grism on 4 August 2023; the black line shows the spectrum in the G140H/100LP grism on 23 August 2023. The uncertainty (rms) is indicated for both profiles. We note that the wiggles in the profiles are in part due to the rotational transitions in the spectra (e.g., see Figure 17, and de Kleer, de Pater, and Ádámkovics (2019)). (b) Two forbidden [SI] lines at 1.0824 and 1.1308 µm, both obtained in the G140H/100LP grism on 23 August 2023. The uncertainty is indicated; because it is very small, the size of the symbol is equal to 3× the rms.</figcaption></figure>

## 4.1. SO Emission Maps and Spectra

Images of the SO emissions are shown in Figure 14. These are constructed from the G140H image data cubes after the thermal hot spot contributions were subtracted, as discussed above. Three images are shown: one averaged over the entire emission region (panel b), one averaged over just the center of the band (panel a; 1.705–1.709 µm; see Figure 13a), and one averaged over the short-wavelength wing of the band (panel c; 1.685–1.703 µm). The latter contains the 1.69 µm peak of the SO emissions, which has been attributed to emissions that are not in Local Thermodynamic Equilibrium (LTE) (i.e., non-LTE) (Bernath & Bittner, 2020; de Kleer, de Pater, & Ádámkovics, 2019; de Pater, de Kleer, & Ádámkovics, 2020; de Pater et al., 2002).

A reprojection of these three images is shown in Figure 15, with panel a showing the thermal continuum emission from Figure 5b (this time on a linear scale). On this panel we have indicated several volcanoes that will be referred to in the discussions below. As in November 2022, we see clear SO emissions above Kanehekili Fluctus. At the center of the line (panel c), the emission is located at the peak intensity of its thermal hot spot emission, but it is slightly elongated in the north-south direction. An almost equally bright region is seen to the north, at or slightly east of the Ukko and Gauwa Paterae. Reprojections of the SO center-band image for both November 2022 and August 2023 are also shown in panels (e) and (f) of Figure 9.

Figure 15b shows the emission averaged over the entire SO band, and here we see that the emission in the north is more intense than that over Kanehekili. It also has a slightly larger longitudinal extent compared to that in just the

SO on Io with JWST/NIRSPEC on 23 August 2023

<figure><img src="image_16.png" /><figcaption>**Figure 14.** Images of the SO emissions: (a) Image averaged over the center of the band (18 channels). Peak intensity is $2.4 \times 10^{11}$ photons s$^{-1}$ cm$^{-2}$ µm$^{-1}$ sr$^{-1}$. (b) Image averaged over the entire band (151 channels). Peak intensity is $9.3 \times 10^{10}$ photons s$^{-1}$ cm$^{-2}$ µm$^{-1}$ sr$^{-1}$. (c) Image averaged over the short wavelength wing of the band (75 channels), that includes the non-LTE 1.69 µm peak. Peak intensity is $9.4 \times 10^{10}$ photons s$^{-1}$ cm$^{-2}$ µm$^{-1}$ sr$^{-1}$. The wavelength region over which the images were averaged is indicated on each panel. The circle outlines Io's limb, and the small filled circle in the lower right shows the approximate resolution of the data (~0.2"—Figure 3). Io North is pointed up. The intensity scale is linear, and the colors are tweaked such that both bright and faint emissions are visible.</figcaption></figure>
<figure><img src="image_17.png" /><figcaption>**Figure 15.** Reprojections of the SO images in Figure 14 superposed on the same USGS map of the satellite as in Figure 5. All maps are normalized such that the maximum intensity is set equal to 100. (a) Intensity at the center of the SO band (same as Figure 5b, i.e., thermal + SO emissions). Several volcanoes are indicated on this panel. (b) Intensity of the SO emissions averaged over the entire SO band. (c) Intensity of the SO emissions averaged over the center of the band. (d) Intensity of the SO emissions averaged over the short wavelength wing of the emission band. The wavelength range over which the images were averaged is indicated on each panel.</figcaption></figure>

center of the band; it extends all the way from Ukko Patera to Fuchi Patera, that is, roughly 60° in longitude. The short-wavelength wing emissions (panel c) show a peak over Fuchi Patera, and one just east of the Ukko and Gauwa Paterae. The strongest non-LTE emissions, though, originate east of Kanehekili Fluctus, just south of Karei Patera, with a second, much fainter, peak near Janus Patera.

As summarized by Perry et al. (2025), *Galileo* SSI detected a plume originating in the Kanehekili Fluctus/Guaraci Fluctus region in 1997, and *Juno/JunoCam* detected two plumes at Kanehekili Fluctus in February 2024: one in the south at 20.7°S, 31.5°W, and one in the north near 15°S, 32°W (see Ravine et al., 2024), which correspond well with the elongated structure of the SO emissions in the center of the band (Figure 15c). JIRAM has consistently observed thermal emissions from Gauwa Patera, and not much from Ukko Patera. The Gauwa Patera was also active in 2023 December 20 Keck data (Figure 4). Hence the SO emissions in the north-west might be related to Gauwa Patera.

In Figure 16 we show spectra at different locations on Io's disk. As before, these spectra were obtained with IDL's program *apphot.pro*, with a radius of 0.05" after the size of each image frame in the data cube was increased by a factor of 5 to create sub-spaxels at a 0.02"-sized grid. Note, though, that the various spectra are not always independent of each other, since the PSF, including its wings, is large. In addition, spectra at locations where the continuum flux is high suffer from a ripple pattern (e.g., spectra 1 and 11), due to the undersampling of the PSF, even though this emission was removed before creating the spectra. Nevertheless, as shown, spectra do appear to vary across the disk, in particular the intensity of the short-wavelength wing relative to the main peak, as was already evident in Figures 14 and 15, and noticed before in Keck data (de Pater, de Kleer, & Ádámkovics, 2020). The 1.69 µm peak sometimes appears to be double (spectra 3–5); if this is real, it may contain interesting clues as to the origin of the 1.69 µm feature.
## SO spectra at different locations on Io's disk

<figure><img src="image_18.png" /><figcaption>**Figure 16.** SO spectra at various locations on Io's disk, each averaged over an area with a radius of 0.5 pixel (0.05"; or 2.5 pixels in the expanded views shown in the images). The center image is from Figure 14b; the small filled circle in the lower right shows the approximate resolution of the data. Spectrum 1 is taken at the peak of Kanehekili's thermal emission (Figure 3). The x-axis of each spectrum covers 1.68–1.73 µm, and the y-axis is from –5 × 10¹⁰ up to 4 × 10¹¹ photons s⁻¹ cm⁻² µm⁻¹ sr⁻¹. These spectra were produced straight from the data cubes (i.e., no corrections were applied).</figcaption></figure>

### 4.2. SO Emission Rates and Column Densities

Figure 13a showed a spectrum of the forbidden SO $a^1Δ \rightarrow X^3Σ^−$ rovibronic transition at 1.707 µm. These emissions, observed when Io was in eclipse, are spin-forbidden and will have electron excitation cross-sections that peak at low energy, ~10 eV. In order to derive the emission rate $F(\text{SO})$ in photons s⁻¹ and column densities N(SO) in cm⁻² we integrate the flux density ($I$ [Equation 1] ×$p_{sr}$) over wavelength and Io's disk to obtain the observed flux, $F_{\text{obs}} = 0.0255 \pm 0.0014$ photons s⁻¹ cm⁻² (Section 2.1), where the uncertainty is based on the standard deviation in the spectrum multiplied by $\sqrt{N}$ (N is the number of spectral points), together with a 5% absolute calibration error (https://jwst-docs.stsci.edu/), added in quadrature. We can convert this to a distance-independent photon emission rate, $F(\text{SO}) = 4πΔ^2F_{\text{obs}}$, with the Io-JWST distance Δ in cm. The total emission
<table><caption>**Table 7**
Constants for SO and [SI] Calculations</caption><thead><tr><th>Transition</th><th>Wavelength<br>μm</th><th>ΔE<br>eV</th><th>Einstein Coefficient<sup>a</sup><br>s<sup>-1</sup></th><th>Lifetime<br>s</th><th>g<sub>u</sub>(upper)</th><th>g<sub>t</sub>(lower)</th></tr></thead><tbody><tr><td>SO <i>a</i><sup>1</sup>Δ → <i>X</i><sup>3</sup>Σ<sup>-</sup></td><td>1.707</td><td>0.727</td><td>0.067</td><td>14.9</td><td>2</td><td>3</td></tr><tr><td>SI <sup>1</sup>D → <sup>3</sup>P<sub>2</sub></td><td>1.0824</td><td>1.14703</td><td>0.0213</td><td>47</td><td>5</td><td>5</td></tr><tr><td>SI <sup>1</sup>D → <sup>3</sup>P<sub>1</sub></td><td>1.1308</td><td>1.0978</td><td>0.0062</td><td>161</td><td>5</td><td>3</td></tr><tr><td>SI <sup>1</sup>D → <sup>3</sup>P<sub>0</sub></td><td>1.154</td><td>1.076</td><td>2.78 × 10<sup>-6</sup></td><td>4.2 days</td><td>5</td><td>1</td></tr><tr><td>SI <sup>1</sup>S → <sup>1</sup>D</td><td>0.7725</td><td>1.55</td><td>1.3</td><td>0.77</td><td></td><td></td></tr></tbody></table>

rate in August 2023 is $F(SO) = (1.53 \pm 0.09) \times 10^{27}$ photons s$^{-1}$. This number is smaller than the emission rate in 1999 (de Pater et al., 2002) and November 2022 (both $2.3 \times 10^{27}$ photons s$^{-1}$), but agrees well with the SO emission rates published by Laver et al. (2007) and de Kleer, de Pater, and Ádámkovics (2019), which range from $(0.8 \pm 0.2) \times 10^{27}$ to $(2.5 \pm 0.6) \times 10^{27}$ photons s$^{-1}$.

If we assume that the observed flux, $F_{obs} = 0.0255 \pm 0.0014$ photons s$^{-1}$ cm$^{-2}$, is emitted from Io's entire surface, we can divide $F_{obs}$ by $\Omega$, the solid angle subtended by Io's 1.094'' disk, and multiply by $4\pi$ to get the $4\pi I$ brightness of $(1.45 \pm 0.09) \times 10^{10}$ photons s$^{-1}$ cm$^{-2}$, or 14.5 kR. Dividing this number by the transition probability (i.e., Einstein coefficient) of $0.067$ s$^{-1}$ (Bernath et al., 2022), gives a column density of SO molecules in the upper state $N(SO[a^1\Delta]) \sim 2.2 \times 10^{11}$ cm$^{-2}$ (see Tables 7 and 8 for useful constants and a summary of results). We note that the Einstein coefficient was recently revised from $2.2 \times 10^{-1}$ (Klotz et al., 1984) to $0.067 \times 10^{-1}$, that is, decreased by a factor of 30 (Bernath et al., 2022).

Although the above calculations assumed the SO emissions to be emitted equally across Io's surface, our images clearly show that SO is spatially distributed in a quite non-homogeneous way. Hence, the column density derived above is too low for areas where SO has actually been detected. It may therefore be more appropriate to consider specific regions on Io, that is the high emission areas we see on Figure 14, referred to as SO spot NW, NE, and Kanehekili, as indicated on Figure 17. These spectra are averaged over a circular area as indicated on the insets, with radii of 0.1'' for the northern spots and 0.2'' for Kanehekili (The spectra are discussed in more detail in Section 4.3.).

Assuming the emissions from these areas to arise from unresolved vents on Io's surface, as expected for volcanic eruptions, we determine (Section 2.1) the $4\pi I$ intensity of Spots NW, NE, Kanehekili to be $9.4 \times 10^{10}$, $8.5 \times 10^{10}$, and $3.3 \times 10^{11}$ photons s$^{-1}$ cm$^{-2}$, respectively. Accounting further for the missing flux due to the PSF (Section 3.2) will increase the intensities in the northern hemisphere by a factor of 2.5, and for Kanehekili by 1.25. Dividing the resulting numbers by the transition probability of $0.067$ gives us the column densities $N(SO[a^1\Delta])$ of $3.5 \times 10^{12}$, $3.2 \times 10^{12}$, and $6.2 \times 10^{12}$ cm$^{-2}$, resp.

While interplanetary Ly-$\alpha$ could in principle produce SO emission through dissociation of $SO_2$, it is insufficient to account for the observed SO brightness as discussed in Section 4.3 of de Pater, Lellouch, et al. (2023). Furthermore, in eclipse the molecular atmosphere of $SO_2$ collapses and SO reacts on the surface with itself to produce the dimer ($S_2O_2$) which in turn reacts with SO to form disulfur monoxide $S_2O + SO_2$, decreasing the SO abundance even more. $S_2O$, however, is not detected in our spectra (Section 6).

By analogy with isoelectronic $O_2(a^1\Delta)$, de Pater et al. (2002) estimated that torus electron collisions with SO ($X^3\Sigma^-$) to excite the molecules to the SO ($a^1\Delta$) state would require a total SO column density of $\sim 5 \times 10^{17}$ cm$^{-2}$ to yield the observed SO 1.707 µm emission rate, a column density in substantial excess of column densities observed in direct sunlight (de Pater, Goldstein, & Lellouch, 2023; Lellouch et al., 2007).

Thus, the torus electrons, whose ultimate source of energy is Jupiter's rotation, are insufficient to explain the SO band brightness and we are left to invoke the tidal energy source that powers Io's volcanic eruptions (Peale

<footnote><sup>a</sup>Einstein coefficients for SO and S were taken from Bernath et al. (2022) and the National Institute of Standards and Technology (NIST) Podobedova et al. (2009), resp. Note, though, that NIST assigns accuracy grades of E for the 1.0824 µm line and D for the 1.1308 µm line, though the observed brightness ratio of these two lines is consistent with the Einstein coefficients.</footnote>
<figure><img src="image_19.png" /><figcaption>**Figure 17.** SO spectra with superposed models (all scaled to the maximum observed intensity) at various locations on Io's disk: (a) Radiance averaged over the disk, (b) Radiance over Kanehekili Fluctus, averaged over a circular area with a radius of 2 pixels (0.2"). (There is an artifact near 1.68 µm.) (c) Radiance over the bright spot North-West and (d) North-East, averaged over an area with a radius of 1 pixel (0.1"). The following models are superposed: (a) Red: Two-temperature model, where the total intensity is: $I = I(200\,\text{K}) + f \times I(1,500\,\text{K})$, with $f = 0.25$. Cyan: Two-temperature model with $I = I(300\,\text{K}) + f \times I(2,400\,\text{K})$, with $f = 0.10$. (b) Red: 2T model, where: $I = I(200\,\text{K}) + f \times I(1600\,\text{K})$, with $f = 0.25$. (c, d) Red: 3T model, where: $I = I(200\,\text{K}) + f_1 \times I(1,500\,\text{K}) + f_2 \times I(2,400\,\text{K}, J > 60)$, with $f_1 = 0.4; f_2 = 0.5$ at $J > 60$ and $\lambda < 1.70\,\mu\text{m}; f_2 = 0.0$ at $\lambda > 1.70\,\mu\text{m}$. Uncertainties (rms) are indicated on each plot. Note that some wiggles in the spectra may be real, as shown by the model spectra.</figcaption></figure>

et al., 1979). In the following we determine thermodynamic equilibrium distributions of the energy levels of SO as a function of temperature to evaluate the potential of a volcanic source (see also de Pater et al. (2002)).

In sunlight, right before entering or after exiting eclipse, Io's atmosphere has SO₂ column densities of ~1 – 2 × 10¹⁶ cm⁻² (de Pater, Luszcz-Cook et al., 2020); with the SO mixing ratio in the range of 3%–10% (de Pater, Goldstein, & Lellouch, 2023; Lellouch et al., 2007), we expect an SO abundance of order 0.3–2 × 10¹⁵ cm⁻². We observed a SO column density in the upper state of N(SO[a¹Δ]) ≈ 2.2 × 10¹¹ cm⁻² using the disk-averaged values, and N(SO[a¹Δ]) ≈ 3.2 – 6.2 × 10¹² cm⁻² for the regions in the northern hemisphere and Kanehekili. Assuming thermodynamic equilibrium, Boltzmann's equation provides the ratio of SO in the upper to the lower state:
<table><caption>**Table 8**
SO Quantities as Observed and Modeled Based Upon Thermodynamic Equilibrium</caption><thead><tr><th></th><th>Δ<br>AU</th><th>N(SO[a<sup>1</sup>Δ])<sub>global</sub><br>cm<sup>-2</sup></th><th>N(SO[a<sup>1</sup>Δ])<sub>N.spots</sub><br>cm<sup>-2</sup></th><th>N(SO[a<sup>1</sup>Δ])<sub>Kan.</sub><br>cm<sup>-2</sup></th><th>F(SO)[a<sup>1</sup>Δ]<br>photons s<sup>-1</sup></th></tr></thead><tbody><tr><td>Observed 2023</td><td>4.611</td><td>(2.2 ± 0.13) × 10<sup>11</sup></td><td>(3.3 ± 0.2) × 10<sup>12</sup></td><td>(6.2 ± 0.6) × 10<sup>12</sup></td><td>(1.5 ± 0.1) × 10<sup>27</sup></td></tr><tr><td>Observed 2022</td><td>4.308</td><td>(3.3 ± 0.5) × 10<sup>11</sup></td><td>X</td><td>X</td><td>(2.3 ± 0.4) × 10<sup>27</sup></td></tr><tr><td>Observed 1999</td><td>4.08</td><td>(3.2 ± 0.6) × 10<sup>11</sup></td><td>X</td><td>X</td><td>(2.3 ± 0.35) × 10<sup>27</sup></td></tr></tbody></table><table><thead><tr><th></th><th>Temperature<br>(K)</th><th>N(SO) (global)<br>cm<sup>-2</sup></th><th>N(SO) (N.spots)<br>cm<sup>-2</sup></th><th>N(SO) (Kan.)<br>cm<sup>-2</sup></th><th>F(SO)[a<sup>1</sup>Δ]<br>photons s<sup>-1</sup></th></tr></thead><tbody><tr><td>Model</td><td>1,000</td><td>1.6 × 10<sup>15</sup></td><td>2.3 × 10<sup>16</sup></td><td>4.3 × 10<sup>16</sup></td><td>7.2 × 10<sup>25</sup></td></tr><tr><td>Model</td><td>1,300</td><td>2.2 × 10<sup>14</sup></td><td>3.4 × 10<sup>15</sup></td><td>6.4 × 10<sup>15</sup></td><td>5.0 × 10<sup>26</sup></td></tr><tr><td>Model</td><td>1,500</td><td>9.3 × 10<sup>13</sup></td><td>1.4 × 10<sup>15</sup></td><td>2.6 × 10<sup>15</sup></td><td>1.2 × 10<sup>27</sup></td></tr><tr><td>Model</td><td>1,700</td><td>4.8 × 10<sup>13</sup></td><td>7.2 × 10<sup>14</sup></td><td>1.4 × 10<sup>15</sup></td><td>2.3 × 10<sup>27</sup></td></tr><tr><td>Model</td><td>2,000</td><td>2.3 × 10<sup>13</sup></td><td>3.4 × 10<sup>14</sup></td><td>6.4 × 10<sup>14</sup></td><td>4.9 × 10<sup>27</sup></td></tr><tr><td>Model</td><td>2,500</td><td>9.8 × 10<sup>12</sup></td><td>1.5 × 10<sup>14</sup></td><td>2.8 × 10<sup>14</sup></td><td>1.1 × 10<sup>28</sup></td></tr><tr><td>Model</td><td>3,000</td><td>5.6 × 10<sup>12</sup></td><td>8.4 × 10<sup>13</sup></td><td>1.6 × 10<sup>14</sup></td><td>2.0 × 10<sup>28</sup></td></tr></tbody></table>

$$ \frac{\mathrm{N}(\mathrm{SO}[a^1\Delta])}{\mathrm{N}(\mathrm{SO}[X^3\Sigma^-])} = \frac{g_u}{g_l} \exp(-\Delta E/T), \qquad (5) $$

with $\Delta E = 0.727$ eV, statistical weights $g_u$ and $g_l$ from Table 7, and the temperature $T$ in eV. Hence, in thermodynamic equilibrium, the total SO column density depends on the temperature of the (hot) volcanic gas, and can be calculated with Equation 5. For example, at 1,300 K (equal to 1300/11605 = 0.112 eV), given an observed N(SO[a^1Δ]) column density of $2.2 \times 10^{11}$ cm⁻², the total SO column density in August 2023 is $2.2 \times 10^{14}$ cm⁻²; for the northern emission regions, this is $3.4 \times 10^{15}$ cm⁻², and $6.4 \times 10^{15}$ cm⁻¹ for Kanehekili. From the various examples shown in the bottom half of Table 8, and the fact that the N(SO) column density is expected to be of order $0.3-2 \times 10^{15}$ cm⁻², it is clear that the SO above the northern spots and Kanehekili must have been quenched (i.e., equilibrated) at temperatures of order 1,500–2,000 K, either in the vent/conduit or the plume. If equilibration occurred in the conduit, these temperatures might reflect magma eruption temperatures, and imply a composition in the upper-basaltic to basaltic-komatiitic composition range.

We can also calculate the expected F(SO) emission rates assuming thermodynamic equilibrium considerations, using:

$$ F(\mathrm{SO}[a^1\Delta]) = r_{[\mathrm{SO}/\mathrm{SO}_2]} \times F(\mathrm{SO}_2) \frac{2}{3} \exp(-\Delta E/T). \qquad (6) $$

Assuming a total SO₂ flux of F(SO₂) = $5 \times 10^{30}$ molecules s⁻¹ (Strobel & Wolven, 2001) and SO/SO₂ ratio $r_{[SO/SO_2]} = 0.10$, we calculate the emission rates shown in the last column of Table 8. A comparison of the modeled and observed flux rates also point at SO quenching temperatures of order 1,500–1,700 K.

### 4.3. SO Spectral Shape and Spatial Distribution

Figures 14–16 show both the spatial and spectral distribution of the forbidden SO $a^1\Delta \rightarrow X^3\Sigma^-$ rovibronic transition at 1.707 µm. The spatial distribution in the short-wavelength wing of the emission feature is quite different from that at the center of the emission band, which is also imprinted on the emission averaged over the entire band. As mentioned above, when we compare the spectra in Figure 16, we see variations in the 1.69 µm emission bump relative to the main peak across Io's disk. This emission bump is most visible in the northern hemisphere (panels 3–6). Most of the emissions appear to originate near Ukko/Gauwa Patera and Fuchi Patera.
<figure><img src="image_20.png" /><figcaption>**Figure 18.** The various transition branches of SO at three different rotational temperatures (200, 500, and 1,800 K). Note the relative changes that take place when the temperature is increased.</figcaption></figure>

We show the disk-averaged spectrum in Figure 17a, together with 3 representative profiles at different locations: at the center of Kanehekili Fluctus (panel b), and at the maximum intensities in a SO-bright spot in the NW (panel c), and in the NE (panel d), as indicated on the insets in the figures. In panel a we superpose several models (all models are normalized to the observed peak intensity): one model (red) shows essentially the 2-temperature fit from de Kleer, de Pater, and Ádámkovics (2019), which was derived from a fit to high spectral resolution data, with the total intensity: *I*(200 K) + *f* × *I*(1,500 K), where the 200 and 1,500 K correspond to the rotational temperatures (Trot) of the gas. In de Kleer, de Pater, and Ádámkovics (2019) we used the SO linelist of Setzer et al. (1999) with *f* = 5/6. In the present paper we use the most recent line list from Bernath et al. (2022), which has rotational transitions up to *J* = 75. With this linelist, *f* = 0.25, which is the same factor as derived by de Pater, de Kleer, and Ádámkovics (2020) when using an earlier version of Bernath's linelist (Bernath & Bittner, 2020). We only use the 0-0 band, since features from the 1-1 band near 1.717 µm are not present in the data (see also Figure 17 in de Pater, de Kleer, and Ádámkovics (2020)).

When increasing the rotational temperature, the profile gets broader at the center, and the intensity in the wings at 1.69–1.70 and 1.71–1.73 µm increases, as exemplified, for example, in Figures 14 and 17 in de Pater, de Kleer, and Ádámkovics (2020). The broadening and the changes in the relative intensities are quite obvious when comparing the transition branches at different temperatures in Figure 18. The red line profile in Figure 17a is almost the same as the cyan profile, modeled as: *I*(300 K) + *f* × *I*(2,400 K), with *f* = 0.10. However, none of these profiles match the 1.69 µm bump.

In Figure 17b we show the emission above Kanehekili Fluctus. As shown, here we do not see emission at 1.69 µm, and the profile can be matched reasonably well with a superposition of a 200 and 1,600 K model. We note that this spectrum is affected more than the other profiles by the wavelength-dependent wiggles due to the undersampling of the PSF, since Kanehekili Fluctus is very bright in the continuum (Section 3).

In Panels (c) and (d) we show spectra of the SO-bright spots in the northern hemisphere. These regions do not show strong continuum emissions, yet show very clear SO emissions, as well as a strong 1.69 µm emission bump. The model superposed on the data consists of a 3-temperature model: *I* = *I*(200 K) + *f*₁ × *I*(1,500 K) + *f*₂ × *I*(2,400 K, at *J* > 60), with *f*₁ = 0.4. We only use high-J values (*J* > 60) at Trot = 2,400 K, with *f*₂ = 0.5 at wavelengths *λ* < 1.70 µm, and we ignore (*f*₂ = 0.0) the high-J values at *λ* > 1.70 µm. Such a combined profile provides a reasonable (by eye) match to the observed spectra; the double-peak at 1.69 µm is modeled by the Trot = 2,400 K curve at *J* > 60 in the rR12-rQ1 branches. This same feature in the high-J values in the qR13-qQ12 branches is not seen in the data, and we therefore used *f*₂ = 0 in the model. The data clearly hint at quite a selective non-LTE effect to give rise to the 1.69 emission bump (see also de Pater, de Kleer, and Ádámkovics (2020)); we are not aware, however, of a physical mechanism that could produce this, although at the high temperatures required to explain the data, perhaps odd non-LTE effects may be introduced by radiative/dissociative recombination of SO⁺/SO₂²⁺ and charge exchange.
<figure><img src="image_21.png" /><figcaption>Sulfur on Io with JWST/NIRSPEC on 23 August 2023</figcaption></figure>

**Figure 19.** Sulfur emissions on Io's disk. (a) Image integrated over 3 channels centered at the center of the [SI] 1.08 µm band (1.08221–1.08268 µm). Peak intensity is $2.6 \times 10^{12}$ photons s$^{-1}$ cm$^{-2}$ µm$^{-1}$ sr$^{-1}$. (b) Image integrated over 3 channels centered at the center of the [SI] 1.13 µm band (1.13062–1.13109 µm). Peak intensity is $7.9 \times 10^{11}$ photons s$^{-1}$ cm$^{-2}$ µm$^{-1}$ sr$^{-1}$. The circle outlines Io's limb, and the small filled circle in the lower right shows the approximate resolution of the data (~0.2"). Io North is pointed up. The intensity scale is linear, and the colors are tweaked such (by changing the slope and minimum values) that both bright and faint emissions on panels (a) and (b) are visible. The angularity (right angles) in the emissions are artifacts, caused by the interpolation on a finer grid. (c) As panel (a), with superposed the projections of the magnetic equator of Jupiter's magnetic field (solid thick white line), the magnetic field direction projected on the plane of the disk (blue arrow), and plasma flow (green arrow). The direction of the magnetic field out of the plane has a significant contribution as displayed in Figure 24. A latitude/longitude grid is overlain, where the sub-jovian meridian (0°W. longitude) is indicated by the thick white dashed line.

**5. Emission Line Maps and Spectra of Atomic Sulfur**

Images of the sulfur emissions, extracted from the G140H image data cubes after the thermal hot spot contri-
butions were subtracted as discussed in Section 4.1, are shown in Figures 19a and 19b. Images are centered on
each of the two lines, 1.082 and 1.1308 µm. A reprojection of these images, superposed on the USGS map, is
shown in Figure 20.

The sulfur emissions are clearly concentrated in the northern hemisphere, and cover a similar range in longitude (~60°, from roughly Ukko Patera through Fuchi Patera) and latitude (~25–40°N) as the SO emissions in the north, although there are a few significant differences. In addition to the fact that no emissions are evident from the area around Kanehekili Fluctus, the S emissions in the northern hemisphere are quite “smooth”, not as “patchy” as the SO emissions. Since these images are part of the same image data cube as the one from which the SO images were produced, the difference cannot be explained by instrumental effects.

<figure><img src="image_22.png" /></figure>

**Figure 20.** Reprojections of the sulfur images in Figure 19 superposed on the same USGS map of the satellite as in Figures 5 and 15. The maps are normalized such that the maximum intensity is set equal to 100. (a) Intensity at 1.08 µm. (b) Intensity at 1.13 µm.
Neutral sulfur lines were first detected with the International Ultraviolet Explorer (Ballester et al., 1987), and later mapped with Hubble Space Telescope's (HST) Imaging Spectrograph (STIS) and Faint Object Spectrograph (Feaga et al., 2002; Roesler et al., 1999). All these observations were carried out in the ultraviolet, between roughly 115 and 230 nm. The images typically show equatorial auroral spots oscillating up and down about Io's equator caused by the threading of Jupiter's magnetic field lines through Io's atmosphere (Saur et al., 2000). The spots are located on the sub- and anti-jovian sides. The brightness decreases when the satellite enters Jupiter's shadow (Roth et al., 2014). In addition to the bright equatorial spots, UV images display a faint limb glow and some extended emission (Wolven et al., 2001).

Io was observed by the *Cassini* spacecraft in late 2000 and early 2001 while the spacecraft was en-route to Saturn. Filters at UV, visible and near-IR wavelengths were used to observe Io while in eclipse. Geissler et al. (2004) used models from Saur et al. (1999); Saur et al. (2002, 2003) which describe the electrodynamic interaction between Io and the plasma torus (Section 5.4) to determine column abundances and mixing ratios of various constituents in Io's atmosphere that would match the *Cassini* multi-filter observations for the anti-jovian spot. The mixing ratios of SO₂, S₂, and O were determined to be roughly 80%, 10%, and 5%, resp. About 2.3 kR was attributed to [SI] at a 2.5% mixing ratio, emitting at 0.7725 µm. This would correspond to a column density of ~1.25 × 10¹⁵ cm⁻² for a SO₂ column density of 5 × 10¹⁶ cm⁻². We note, though, that analysis of these data required difficult estimated background subtraction methods, and most of the total observed emission (~20 kR) was allocated to the [KI] D lines at 766.5 and 769.9 nm that dominated the IR1 filter intensity.

At first glance, the morphology of the sulfur emissions observed with JWST in this paper is quite different from that observed with *Cassini* and HST. The apparent disk-averaged brightnesses ($4\pi I$ as for SO in Section 4.2) of the 1.082 and 1.1308 µm transitions are 4.2 and 1.4 kR, respectively, which are similar to the brightness of the 0.7725 µm line mentioned above. The main apparent difference between the emissions may be the morphology, since the JWST emissions are confined to the northern hemisphere on the disk, in contrast to the UV- and visible-light detected auroral spots at Io's limb. We return to this point in Sections 5.4 and 5.5, where we discuss the role of the viewing geometry and the physical processes causing the emissions. Since the *Cassini* observations were measured for just the anti-jovian spot, and the JWST emissions are clearly disk-resolved and confined to a strip in the northern hemisphere, it is more appropriate to use the average brightness of the most intense emissions in the northern hemisphere. These brightnesses ($4\pi I$) are 7.4 and 2.25 kR at 1.082 and 1.1308 µm, respectively.

In Figure 21 we show spectra at different locations on Io's disk. As expected from the image at the center of the frame, the intensity of the spectral peaks varies somewhat across the disk. Note, though, that the various spectra are not always independent from each other, since the PSF, including its wings, is large.

## 5.1. [SI] 1.082 and 1.1308 µm Emission Rates and Local Thermodynamic Equilibrium (LTE)

The two detected sulfur lines, [SI] 1.082 and 1.1308 µm, correspond to spin-forbidden transitions of the S(¹D) to the ground multiplet S(³P₂,₁) fine structure levels. A Grotrian diagram is shown in Figure 22.

The three spin-forbidden transitions from the S(¹D) to the ground multiplet S(³P₂,₁,₀) have the same upper energy level (1.147 eV) and the probability that one photon is emitted at 1.154 µm is equal to the ratio of its Einstein coefficient divided by the sum of the coefficients for all three transitions (Table 7), $1 \times 10^{-4}$. This corresponds to a brightness of ~1 R, much smaller than the 3-σ upper limit in our disk-averaged data of 60 R for a typical [SI] ~1-µm line. It is therefore not surprising that the 1.154 µm line is not detected in our data.

The total wavelength and disk-integrated [SI] flux, $F_{obs}$, is equal to 0.0073 and 0.0025 photons s⁻¹ cm⁻² at 1.082 and 1.1308 µm, respectively. If we multiply these by $4\pi\Delta^2$ (Section 2.1), we get a total photon emission rate of $5.9 \times 10^{26}$ photons s⁻¹ from the upper, S(¹D), state. The disk-averaged $4\pi I$ brightness of the two lines is 4.2 and 1.4 kR, respectively.

However, like for the SO spatial distribution, it would make more sense to consider the apparent $4\pi I$ brightnesses of these lines in just the northern hemisphere, which are 7.4 and 2.25 kR for the 1.08 and 1.13 µm lines, respectively. Note that the ratio of these numbers is 3.3, very close to the ratio of the transition probabilities, 0.0213/0.0062 = 3.43, as expected.
S spectra at different locations on Io's disk

<figure><img src="image_23.png" /></figure>

**Figure 21.** Sulfur spectra at various locations on Io's disk, each averaged over an area with a radius of 0.5 pixel (or 2.5 pixels in the expanded views shown in the images). The center image is from Figure 19b; the small filled circle in the lower right shows the approximate resolution of the data. Spectrum 1 is taken at the peak of Kanehekili's thermal emission (Figure 3). The horizontal axis of each spectrum covers 1.065–1.15 µm, and the vertical axis is from –10¹¹ up to 1.3 × 10¹² photons s⁻¹ cm⁻² µm⁻¹ sr⁻¹. As in Figure 16, these sulfur spectra were produced straight from the data cubes (i.e., no corrections were applied). Ratios R are indicated on the figure: R_S is the ratio between the [SI] intensities at 1.08 and 1.13 µm (see also Figure 19c); R SO is the ratio between the [SI] intensity at 1.08 µm with the SO intensity integrated over the full bandwidth (1.685–1.721 µm). Typical uncertainties in R_S range from ~0.14 to ~0.3 in strong to weak lines. The uncertainties in the ratio R SO are also in the range of ~3%–10%. Note: there are some spurious spikes near 1.145 µm in several panels.

Dividing the observed brightnesses by their respective transition probabilities (Table 7) results in emitting column densities of N(S[¹D]) = 3.47 × 10¹¹ cm⁻² and 3.63 × 10¹¹ cm⁻², respectively. We use the average column density of 3.55 × 10¹¹ cm⁻² below. The globally averaged column density is ~2 × 10¹¹ cm⁻².

Using the statistical weights (Table 7) and Boltzmann's equation, we can write the number density in the various rotational (J) levels:
<figure><img src="image_24.png" /><figcaption>**Figure 22.** Grottrian diagram: Electronic transitions between energy levels of the S atom that give rise to the forbidden 1.082 and 1.131 µm emission lines that are observed with the James Webb Space Telescope, as well as the 1.154 and 0.7725 µm lines.</figcaption></figure>

$$n(^{3}\text{P}_{1}) = n(^{3}\text{P}_{2}) \frac{g(^{3}\text{P}_{1})}{g(^{3}\text{P}_{2})} \exp(-0.049/T), \qquad (7)$$

$$n(^{3}\text{P}_{0}) = n(^{3}\text{P}_{2}) \frac{g(^{3}\text{P}_{0})}{g(^{3}\text{P}_{2})} \exp(-0.071/T), \qquad (8)$$

with *T* in eV.

In the limit of high temperature, the population ratio of these multiplet levels (*J* = 2 : 1 : 0) would be 5:3:1. At 1,700 K, the ratios are 5:2.14:0.62 and the sum of the individual values is 7.76 in comparison to the high temperature limit of 9.

If the total S column density, N(S), in the “ground” (³P) term is determined by LTE thermodynamics, then N(S) can be obtained from:

$$N(S[^{3}\text{P}_{2}]) = 3.55 \times 10^{11} \left[ n(^{3}\text{P}_{2}) \times \frac{g(^{1}\text{D})}{g(^{3}\text{P}_{2})} \exp(-\Delta E(1.082)/T) \right]^{-1}, \quad (9)$$

where $n(^{3}\text{P}_{2}) = 5$, and at $T = 1,700$ K:

$$\left[ n(^{3}\text{P}_{2}) \times \frac{g(^{1}\text{D})}{g(^{3}\text{P}_{2})} \exp(-\Delta E(1.082)/T) \right] = 2 \times 10^{-3}.$$

Hence, at 1,700 K, the total column density N(S) = 3.55 × 10¹¹/(2 × 10⁻³) = 1.8 × 10¹⁴ cm⁻². At 1,500 K, N(S) = 5.1 × 10¹⁴ cm⁻² (Table 9). Assuming an SO₂ abundance of ~1.5 – 5 × 10¹⁶ cm⁻², the S column abundance at 1,500–1,700 K is roughly in agreement with the S/SO₂ abundance ratios observed over Pele (0.3%–0.7% McGrath et al. (2000)). However, they are still roughly a factor of 10 higher than predicted based upon equilibrium chemistry of volcanic gases above Pele (Schaefer & Fegley, 2005). Because of this, and because of the smooth appearance of the S emissions in contrast to SO, we next investigate if collisional excitation by torus electrons could provide a majority, if not all, of the [SI] 1.082 and 1.130 µm lines, even though we did rule out this mechanism for SO.

## 5.2. [SI] 1.082 and 1.1308 µm Emission Rates and Electron Impact Excitation

During the *Voyager* 1 flyby of Jupiter the Plasma Science Experiment (PLS) performed detailed measurements of electron density and temperature distributions. Sittler and Strobel (1987) analyzed this data set in terms of a cold

<caption>**Table 9**</caption>

S Quantities as Observed and Modeled Based Upon Thermodynamic Equilibrium

<table><thead><tr><th></th><th>Δ<br>AU</th><th>N(S)[¹D]<sub>Global</sub><br>cm<sup>-2</sup></th><th>N(S)[¹D]<sub>North</sub><br>cm<sup>-2</sup></th></tr></thead><tbody><tr><td>Observed 2023</td><td>4.611</td><td>(2.1 ± 0.2) × 10<sup>11</sup></td><td>(3.55 ± 0.4) × 10<sup>11</sup></td></tr><tr><th colspan="4">Temperature (K)</th></tr><tr><td>Model</td><td>1,000</td><td>2.5 × 10<sup>16</sup></td><td>4.3 × 10<sup>16</sup></td></tr><tr><td>Model</td><td>1,300</td><td>1.2 × 10<sup>15</sup></td><td>2.0 × 10<sup>15</sup></td></tr><tr><td>Model</td><td>1,500</td><td>3.0 × 10<sup>14</sup></td><td>5.1 × 10<sup>14</sup></td></tr><tr><td>Model</td><td>1,700</td><td>1.0 × 10<sup>14</sup></td><td>1.8 × 10<sup>14</sup></td></tr><tr><td>Model</td><td>2,000</td><td>3.3 × 10<sup>13</sup></td><td>5.5 × 10<sup>13</sup></td></tr><tr><td>Model</td><td>2,500</td><td>9.5 × 10<sup>12</sup></td><td>1.6 × 10<sup>13</sup></td></tr><tr><td>Model</td><td>3,000</td><td>3.5 × 10<sup>12</sup></td><td>6.0 × 10<sup>12</sup></td></tr></tbody></table>
Maxwellian component and a hot super-thermal component, with the former relevant to the [SI] lines. The spacecraft flew through the Io flux tube and when the data were analyzed in terms of an Io-centric coordinate system, the electron temperature decreased from 5.5 eV at 20 R<sub>Io</sub> to 4.8 eV at closest approach (11.5 R<sub>Io</sub>) with a representative electron density of ~2,000 cm<sup>-3</sup>. Measurements by the Galileo/PLS yielded electron temperatures and densities T<sub>e</sub> ~ 4 eV and n<sub>e</sub> ~ 2,000 cm<sup>-3</sup> close to Io (Frank & Paterson, 1999), whereas Gurnett et al. (1996) obtained maximum values of n<sub>e</sub> ~ 4,000 cm<sup>-3</sup> with their plasma wave instrument. Interpretation of Juno radio occultations showed electron densities between roughly 1,500 and 2,500 cm<sup>-3</sup> (Moirano et al., 2021). These measurements suggest that a reasonable bound on the (n<sub>e</sub>, T<sub>e</sub>) parameter space is n<sub>e</sub> ~ 1,500-4,000 cm<sup>-3</sup> and T<sub>e</sub> ~ 4-5 eV.

To calculate the electron impact collisional excitation rate, we adopt the effective collisional strengths computed by Tayal (2004) based on the explicit fine structure of the S(³P) term to the S(¹D) and S(¹S) terms. The collisional strength, Ω<sub>if</sub>(E<sub>ie</sub>/E<sub>if</sub>), is a non-dimensional cross-section related to the actual cross-section by the following equation, where i is the initial level and f is the final level:

$$ \sigma_{if}(E_{ie}) = \frac{\Omega_{if}(E_{ie}/E_{if})}{g_{i}(E_{ie}/27.2 \text{ eV})} \pi a_{o}^{2}, \qquad (10) $$

where a<sub>0</sub> is the Bohr radius, so that πa<sub>0</sub><sup>2</sup> = 8.7974 × 10<sup>-17</sup> cm<sup>2</sup>, and 27.2 eV is equal to 2 Rydbergs = 1 atomic unit. The effective collisional strength is defined as:

$$ \gamma_{if}(T_e) = \int_{0}^{\infty} \Omega_{if} \exp\left(-\frac{E_{if}}{T_e}\right) d\left(\frac{E_{if}}{T_e}\right). \qquad (11) $$

If Ω<sub>ik</sub>(E) is not a function of energy, then γ<sub>if</sub> = Ω<sub>if</sub>. The collisional excitation rate coefficient with E and T<sub>e</sub> in eV is:

$$ \text{Rate}(T_e) = \frac{8.0104 \times 10^{-8}}{g_i T_e^{1/2}} \gamma_{if}(T_e) \exp\left(-\frac{E_{if}}{T_e}\right) \text{cm}^3 \text{s}^{-1}. \qquad (12) $$

With the relevant effective collisional strengths in Table 3 of Tayal (2004), we calculated the excitation rates from the S(³P<sub>1,2,0</sub>) multiplets to the S(¹D) term, as summarized in Appendix C, Table C1.

As discussed above, the typical range of torus electron temperatures at Io's orbit is 4-5 eV which is well matched to the excitation energy of ~1 eV of the 1 µm sulfur lines originating from S(¹D). The next term up is S(¹S) with an excitation energy of ~2.75 eV (Figure 22) and should be very sensitive to the electron temperatures in the torus. Tayal (2004) has also calculated the effective collisional strengths from the S(³P<sub>1,2,0</sub>) multiplets to S(¹S) as a function of electron temperature in eV. From Table 3 of Tayal (2004), we calculated the rates for excitation from the S(³P<sub>1,2,0</sub>) multiplets to S(¹S) (Appendix C, Table C2).

From the S(¹S) level there are two pathways to radiate to a lower term, one to S(¹D) with a photon at 0.7725 µm and transition probability of 1.38 s<sup>-1</sup> and the other to the S(³P<sub>1</sub>) ground state emitting a photon at 458.9 nm with transition probability of 0.303 s<sup>-1</sup>. Thus, the former pathway proceeds ~82% of the time to the ¹D level.

At a specific electron temperature, the sum of the rates from Table C1 (from ³P<sub>1,2,0</sub> → ¹D) and C2 (from ³P<sub>1,2,0</sub> → ¹S), when multiplied by the column densities in each N(S[³P]) multiplet J = 2, 1, 0 state (which are based on an assumed atmospheric temperature) times an assumed electron density yields the apparent brightness in photons cm<sup>-2</sup> s<sup>-1</sup> (or kR after multiplication by 10<sup>-9</sup>). For example, at T<sub>e</sub> = 4 eV with a hypothetical column abundance of ~1 × 10<sup>14</sup> cm<sup>-2</sup> × the sum of the population ratios (Equations 7 and 8) at 1,700 K (=7.76) we get a total sulfur column density N(S) = 8 × 10<sup>14</sup> cm<sup>-2</sup>. After multiplying this with the sum of the rates in the appropriate proportions (5:2.14:0.63 for J = 2, 1, 0 at 1,700 K) and the electron density n<sub>e</sub> = 2,500 cm<sup>-3</sup>, one obtains a (4πl) brightness of 9.65 kR for the combined 1-µm lines. What is important here is the product N(S) × n<sub>e</sub>; for example, N(S) = 5 × 10<sup>14</sup> cm<sup>-2</sup> with n<sub>e</sub> = 4,000 cm<sup>-3</sup> also yields 9.65 kR.
<figure><img src="image_25.png" /><figcaption>**Figure 23.** (a) Total S line emission brightnesses in kR from the S(¹D) term to the S(³P₂,₁) multiplets at 1.082 and 1.130 µm for an assumed electron density *n*<sub>e</sub> = 2,500 cm⁻³ and S emitting column density 8 × 10¹⁴ cm⁻² as a function of *T*<sub>e</sub>, or any other combination of densities as long as the product is 2 × 10¹⁸ cm⁻⁵, when the atmospheric temperature is 1,700 K (0.1465 eV) and the sum of the statistical weights for the S(³P) term is 7.76 (n(³P₂) = 5, n(³P₁) = 2.14, n(³P₀) = 0.62) (solid red line). The green line gives the contribution by direct collisional excitation to the ¹D term. The solid magenta line represents the contribution of collisional excitation first to the ¹S term and then emission of a 0.7725 µm photon to populate the ¹D term and subsequently enhancing the 1.082 and 1.130 µm brightnesses. This line gives thus the brightness of the 0.7725 µm and added brightness to the 1-µm lines, and the red line is the sum of the green and magenta lines. The blue line is the collisional excitation rate (Equation 12; Tayal, 2004) ×10⁹ for the sum of the two 1-µm lines (1.082 + 1.1130 µm) in cm³ s⁻¹. (b) Total S line emission brightnesses in kR from the S(¹D) term in the S(³P₂,₁) multiplets at 1.082 and 1.130 µm as a function of *T*<sub>e</sub>. Calculations are for atmospheric temperatures of 1,700 K (red, as in Panel a), 1,000 K (magenta), 600 K (green), 400 K (blue), and 200 K (brown). The S column density increases from 5.4 × 10¹⁴ cm⁻² at 200 K to 8 × 10¹⁴ cm⁻² at 1,700 K as the relative population densities of the ³P₁ and ³P₀ increase with temperature (Equations 7 and 8). The electron density *n*<sub>e</sub> = 2,500 cm⁻³. The electron temperature in both panels is in eV, and the horizontal black lines are the brightness equal to 9.65 kR.</figcaption></figure>

To illustrate the dependence on electron temperature for direct electron impact excitation, Figure 23a shows a variety of S line emission brightnesses from the S(¹D) to the S(³P) level, assuming N(S) = 8 × 10¹⁴ cm⁻² and *n*<sub>e</sub> = 2,500 cm⁻³. The magenta (solid) line represents the contribution of collisional excitation first to the ¹S term and then emission of a 0.7725 µm photon to populate the ¹D term and subsequently down to the S(³P₂,₁) levels (Figure 22). The green line gives the contribution through direct collisional excitation of S(³P) to the ¹D term. The red (solid) line is the combined effect of these two processes, and corresponds to the 9.65 kR at 4 eV calculated above.

The calculations that yield the 9.65 kR of the 1 µm lines in Figure 23a also yield ~600 R of the 0.7725 µm line for *T*<sub>e</sub> = 4 eV. This is well below the 2.3 kR brightness allocated to sulfur in the *Cassini* eclipse multi-filter observations of the anti-jovian spot (Table 5 of Geissler et al. (2004)), discussed above. However, the [SI] 0.7725 µm transition has recently been measured from Keck/HIRES data, with a brightness well below the *Cassini* value. The Keck/HIRES data were taken in eclipse during ingress (similar to our JWST observations) on four dates in August 2023 and September-October 2024, as well as on 24 November 2022 during egress. On August 25, two days after our JWST observations, the measured disk-averaged brightness was 361 ± 7 R; on all dates together the disk-averaged brightness varies between 325 and 420 R (Milby et al., 2025). Assuming the Keck/HIRES emissions also originate from a narrow latitude range on Io like the JWST 1-µm lines, we should increase the disk-averaged brightness by a factor of ~1.7 (like for the JWST 1-µm lines), which results in 613 R two days after our JWST observations were taken. This agrees extremely well with the 0.7725 µm brightness expected based upon our JWST 1-µm observations.

To show the importance of atmospheric temperature on the calculations, we repeated the calculations in Figure 23a for a series of atmospheric temperatures ranging from 200 to 1,700 K in panel (b). At $T_e = 4$ eV the total [SI] line emission brightness is ~1 kR for the lowest temperature, 200 K; that is, a factor of 10 below the observed total brightness of ~10 kR. It would require an S column density of ~4 × 10¹⁵ cm⁻² to yield the observed total brightness of 10 kR. At an atmospheric temperature of 1,700 K the [SI] line emission brightness does match the observed ~10 kR value for N(S) = 8 × 10¹⁴ cm⁻². These calculations thus show that atmospheric temperature is an important parameter when calculating the [SI] emissions due to electron impact excitation on the ¹S and ¹D states, since the atmospheric temperature determines the population ratios of the various *J* levels of the S(³P) term (see Appendix C). As shown, to excite sulfur from the S(³P) to the S(¹D) and S(¹S) levels, the atmosphere needs to be hot, since it is essential for the S(³P₁,₀) fine structure states of the ground term to be adequately populated as Tables C1 and C2 show that most of the collisional excitation originates from these two states rather than the lowest energy state S(³P₂) (Figure 22). This detail has usually been overlooked in the past (e.g., in Schmidt et al. (2023), Milby et al. (2025)). We note that the older method (Zatsarinny & Tayal, 2002) results in a column density ~3.6 × 10¹⁴ cm⁻² to explain the 1-µm lines, and would predict ~800 R at 0.7725 µm (Appendix C). We further note that a hot atmosphere is expected based upon Alfvénic electrodynamic interactions by torus electrons with Io's highly conducting ionosphere. As shown by Strobel et al. (1994), Joule/resistive heating can elevate the atmospheric temperature to ~1,700 K.
Note that our derived S column density of $8 \times 10^{14} \text{ cm}^{-2}$ may be somewhat higher than the S/SO$_2$ abundance ratio as measured over Pele (McGrath et al., 2000), and at least an order of magnitude higher than predicted to be released from volcanic vents (Schaefer & Fegley, 2005). The latter abundance would have been expected if all S emissions would have originated from sulfur atoms released from the vent in an excited state. The abundance of sulfur atoms in Io's upper atmosphere, however, is expected to be much larger compared to the value based upon equilibrium chemistry due to photodissociation of SO$_2$, SO and S$_2$. Hence a sulfur column density of order $8 \times 10^{14}$ is a very reasonable value.

### 5.3. Dissociative Excitation and Collisional Quenching

While dissociative excitation (through photons or collisional impact) is an additional emission source at some level, there are a few reasons we believe this contribution is negligible as it requires additional excitation energy beyond the 1.15 eV for S(¹D) and 2.7 eV for S(¹S) from S(³P$_2$). The dissociation energy of SO$_2$ is 5.66 eV and for singlet products, SO(¹Δ) + O(¹D), it requires 8.34 eV. The ionization potential of SO$_2$ is 12.349 eV. The dissociation energy of SO is 5.36 eV and for S(¹D) it will be an additional 1.15 eV for a total of 6.51 eV, assuming O remains in its ground ³P$_2$ state. Note that the JWST Io measurements were performed in Jupiter's shadow. The only photons available for dissociative excitation are solar Lyman alpha photons scattered by interplanetary (IPM) H atoms. de Pater, Lellouch, et al. (2023) estimated that for an optically thick SO$_2$ atmosphere these IPM photons would produce $2.3 \times 10^7$ SO(¹Δ) $\text{s}^{-1}$ $\text{cm}^{-2}$ equivalent to 23 R and 23 R of O(¹D) red line at 630 nm. Note, though, that in-eclipse, most of Io's SO$_2$ atmosphere is collapsed (e.g., Summers & Strobel, 1996; de Pater, Luszcz-Cook et al., 2020), and hence the above numbers will be even smaller.

If a non-negligible fraction of the S(¹D) atoms is produced through dissociation of SO, one would expect to see this in the observed morphology of the S 1-µm line emissions. As shown in Figures 14 and 19, there is no hint of the 1-µm emissions being produced through dissociation of SO.

As mentioned in Section 5.2, the best fit to the S(¹D) lines requires a temperature of ~1,700 K, which is consistent with the atmospheric model of Strobel et al. (1994), where substantial Alfvénic heating leads to high temperatures above 250 km. Since sulfur atoms are significantly lighter than SO$_2$ and SO, they can be present at these high altitudes. Dissociations producing S(¹D) at the lower altitudes where SO$_2$ and SO are present are likely collisionally quenched by the background gas in the dense molecular atmosphere, as detailed below.

The lifetime of the S(¹D) term is 36 s based on the Einstein coefficients given in Table 7 for the two observed lines. Collisional quenching rates by SO$_2$ and SO are not available, but there are laboratory measurements and/or theoretical calculations for the inert gases Ar and Kr, and H$_2$ (D. M. G. Williams et al., 2025). The rates are in the range of $1 - 2 \times 10^{-11} \text{ cm}^3 \text{ molecule}^{-1} \text{ s}^{-1}$ at room temperature and below. Thus, quenching densities $\ge 1.4 \times 10^9 \text{ cm}^{-3}$ are required to quench S(¹D) with a radiative lifetime of 36 s. Such a gas density is present in the lower few tens (~40) of kilometers in Io's SO$_2$ atmosphere if Io is in sunlight (Strobel et al., 1994; Walker et al., 2010, 2012), well below the level where we observe the S-atoms (above a few 100 km).

### 5.4. Morphology of [SI] Emission

The [SI] emission shown in Figures 19a and 19b is confined to the northern hemisphere mostly within the disk of Io. In this subsection we discuss how the electron impact mechanism discussed in the previous section might explain the morphology of this emission pattern. Key for this mechanism is the fact that electron energy from the torus is either advected by the plasma flow into Io's atmosphere or conducted along the magnetic field lines into Io's atmosphere (Saur et al., 2000, 2003). Thus plasma flow and magnetic field control the structure of electron temperature in Io's atmosphere.

In Figure 19c we added the magnetic equator to the JWST image of the [SI] emissions, defined as the locations where Jupiter's magnetospheric field is tangent to Io. The background magnetic field is calculated with the JRM33 model (Connerney et al., 2022). The JWST observations were obtained when Io was at $\Psi_m = -9^\circ$ magnetic latitude. Therefore the magnetic equator is rotated toward the north on Io's sub-Jovian hemisphere (see Figures 19c and 24).

As mentioned before, the brightness peaks in the northern hemisphere. This is consistent with the idea that the emission morphology is controlled by the magnetic field and the torus plasma above and below Io, which govern
<figure><img src="image_26.png" /><figcaption>**Figure 24.** Sketch of magnetic field environment when Io is at magnetic latitude of $\Psi_m = -8^\circ$, that is, below the center of the Io torus. Jupiter and James Webb Space Telescope are to the left, the plasma flow is into the plane. Magnetic field lines are tilted toward Jupiter and Io's northern hemisphere is facing the center of the Io torus. Field lines with hot electrons are shown in red and field lines with cold electrons are blue. The bulge around Io's equator represents its atmosphere.</figcaption></figure>

the structure of the electron density and temperature in the ionosphere. The hemisphere of Io which faces the center of the plasma sheet is connected through the magnetic field lines to a larger electron energy reservoir compared to the other hemisphere. On the plasma sheet facing hemisphere, therefore, the electrons are less strongly cooled than on the other hemisphere where the electron energy reservoir is smaller. Therefore electron impact excited emission is expected to be brighter on the hemisphere facing the center of the torus. This effect has been frequently observed within HST observations of Io and Europa (Retherford et al., 2003; Roth et al., 2014, 2016).

The [SI] emission in Figure 19 is mostly north of Io's magnetic equator. The tilt of Jupiter's background magnetic field at the time of the JWST observations contribute to this effect as displayed in Figure 24. In this figure the viewing geometry is along the plasma flow direction, with JWST and Jupiter to the left. The magnetic field lines are strongly tilted toward Jupiter as Io is at $\Psi_m = -9^\circ$. Magnetic field lines with cold electrons facing away from the center of the plasma sheet are shown in blue, while magnetic field lines facing toward the plasma sheet with hot electrons are shown in red. Significant emissions are only generated on field lines with hot electrons. South of the magnetic equator the near surface neutral atmosphere is embedded within cold electrons and little emission is generated. North of the equator the electrons are hotter producing brighter emissions. Additionally, field lines tangent to Io will have hotter electrons further north due to finite electron heat conductivity along field lines. In Figure 24 we additionally assumed that Io's atmosphere during eclipse is confined around its equator and decreases toward the limb with maximum density around that location which was near the subsolar point shortly before Io entered the eclipse.

Of particular relevance for this study is also Figure 25, which shows UV emissions from the brightest lines of the forbidden and allowed transitions of the [SI] 147.9 nm multiplets. Two exposures of these emissions were obtained with HST-STIS on 25 February 2000, and were mapped simultaneously with the [OI] 135.6 nm emission that was shown in Figure 12 of Roth et al. (2014). The [SI] emission was extracted from the combined eclipse exposures following the data processing described in Roth et al. (2014). The data were taken when Io

was at a magnetic latitude $\Psi_m = +8.8^\circ$ and only shows emission in the southern hemisphere, roughly centered around the sub-Jovian longitude (dashed white meridian line). This longitude is to the right of the central meridian in this image, which was taken before egress, in contrast to the JWST image taken after ingress. Localized spot emissions above the limb around the magnetic equator seen by HST at other observing geometries are not visible in the HST and JWST images of Figures 19 and 25. Such localized above-limb spot emissions would be expected in the eclipse images for observations taken from other viewing angles. The sub-Jovian on-disk emission spot/region in Figures 19 and 25 might have an anti-Jovian counterpart that lies behind the disk. Hence, both the JWST and UV images are consistent with the interpretation that the torus facing hemisphere and the direction of the magnetic field control the emission. When Io is outside of eclipse the electron impact excited UV emission is located primarily off the limb as two near-equatorial spots (Roesler et al., 1999; Roth et al., 2014). Unfortunately, it will not be possible to take images of the [SI] 1-µm lines in sunlight due to Io's strong reflected sunlight component.

With S column densities of order $10^{14}-10^{15}$ cm$^{-2}$ there might be concern that the torus electrons would cool down as they use their thermal energy in collisional excitations of atoms and molecules. But while the torus electrons

<figure><img src="image_27.png" /><figcaption>**Figure 25.** Combined UV image of two HST-STIS G140L exposures of Io in eclipse from 25 February 2000. The solid white disk shows Io centered at 147.36 nm between the brightest lines of the forbidden and allowed transitions of [SI] 147.9 nm multiplets. The orange dotted disks indicate the locations of Io at the fainter multiplet lines in the spectral images. For more information on the STIS spectral images the reader is referred to Roth et al. (2014). The dashed meridian in the center shows the sub-jovian longitude. The magnetic equator at the time of the observation is shown in blue.</figcaption></figure>
<figure><img src="image_28.png" /><figcaption>**Figure 26.** Apparent emission rates without allowances for scattering or absorption as brightnesses in kR, with $n_e = 2500 \text{ cm}^{-3}$, S column density = $8 \times 10^{14} \text{ cm}^{-2}$. Red line is total S(1D) 1-µm lines' emission. CIS1474 brown line is the brightness of only the S(³D₀⁻³P₂, 147.41 nm) line and CIS1473 orange is the brightness of only the S(⁵D₀⁻³P₂, 147.3 nm) line in kR, both calculated by the Tayal (2004) method. The Γ1,480 is the multiplet including the 147.41 nm line and Γ1,478 is the multiplet including the 147.3 nm line with respective brightnesses in kR calculated by the method of Zatsarinny and Tayal (2002) for all lines within the respective multiplets.</figcaption></figure>

<figure><img src="image_29.png" /><figcaption>**Figure 27.** The red line and the blue line represent the brightnesses in Rayleighs as a function of altitude (km), respectively, for S[⁵D⁰⁻³P₂] 147.3 nm and S[³D⁰⁻³P₂] 147.41 nm by electron impact with $T_e = 4$ eV (solid lines) and $T_e = 4.3$ eV (dash-dot lines) for the atmosphere above that altitude. The solid lines are for the multiplets calculated with the method of Zatsarinny and Tayal (2002) and the dash-dot lines are for the single lines to the ³P₂ state within the multiplets calculated with the method of Tayal (2004). The magenta line gives the sulfur atom number density (cm⁻³) and the black line the sulfur atom column density (cm⁻²) ×10⁻⁷ as a function of altitude (km).</figcaption></figure>

are losing energy, the accompanying torus ions with $T_i \sim 100$ eV are a collisional reservoir of energy to heat the electrons and maintain their temperature. For example, a test electron at 2 eV inserted in the ionosphere with $2,500 \text{ cm}^{-3}$ electrons at 4 eV and $1,500 \text{ cm}^{-3}$ ions at 80 eV would have a plasma relaxation time constant back to 4 eV of approximately 25 s.

This qualitative explanation needs to be followed up by detailed three-dimensional numerical modeling. Such modeling is however beyond the scope of this manuscript.

### 5.5. Comparison of JWST IR, Keck Visible, and HST UV Emissions

The total brightness of all [SI] 147.9 nm multiplet emissions in Figure 25 within 1.5 R<sub>Io</sub> around the disk center (accounting for multiplet offset and smearing) has been computed and normalized to the area of Io's disk. This brightness is 348 R, which is lower than in all other [SI] images obtained by HST-STIS in sunlight, as shown in Figure 7c in Roth et al. (2014). Like in the IR for JWST, above Io's southern hemisphere we expect a brightness ~1.7 times higher than the disk-averaged value, which is 590 R.

Given that we need a sulfur column density of $8 \times 10^{14} \text{ cm}^{-2}$ to explain the [SI] 1-µm lines in the northern hemisphere (or $4.6 \times 10^{14}$ when normalized to Io's disk) we now turn to see whether we can understand the HST STIS UV lines imaged in Figure 25, which contain the S(⁵D⁰⁻³P₂, 147.3 nm) and S(³D⁰⁻³P₂, 147.41 nm) lines within the observed Γ1,478 and Γ1,480 multiplets, resp., with a combined disk-averaged brightness of 348 R. The latter line is spin allowed with transition probability $A_{ki} = 1.96 \times 10^8 \text{ s}^{-1}$, oscillator strength = 0.0894, statistical weights $g_i = 5$, $g_k = 7$ (Podobedova et al., 2009), and at line center its cross-section is $2.1 \times 10^{-13} \text{ cm}^2$ at 1,700 K and thus optical depth = 1 occurs at a sulfur column density ~$5 \times 10^{12} \text{ cm}^{-2}$, which is a factor of over 100 lower than the column density needed to explain the sulfur 1-µm lines by electron impact. The S(⁵D⁰⁻³P₂, 147.3 nm) line while spin forbidden has a transition probability = $3.73 \times 10^7 \text{ s}^{-1}$, oscillator strength = 0.017, statistical weights $g_i = 5$, $g_k = 7$, and at line center its cross section is $4 \times 10^{-14} \text{ cm}^2$ at 1,700 K and thus optical depth = 1 occurs at a sulfur column density ~$2.5 \times 10^{13} \text{ cm}^{-2}$, also much lower than the column density needed to explain the sulfur 1-µm lines by electron impact.

The UV brightness is computed as before as the integral of the product of the excitation rate times the sulfur density times the electron density with the added factor of the Doppler line shape times the transmission of Doppler lines through the overhead column of sulfur atoms. It is carried out over an effective total line width of 8 Doppler widths as the optical depth at line center for the 147.4 nm line is ~100. (Detailed calculations are provided in Appendix B.) The excitation rates for these two multiplets (Γ1,478 and Γ1,480) were calculated with parameters given in Zatsarinny and Tayal (2002), while we also calculated the excitation rates for the individual 147.3 and 147.41 nm lines terminating in the (³P₂) state with the Tayal (2004) method. The resulting apparent emission rates are shown in Figure 26; note the extreme sensitivity to the electron temperature below 10 eV. The expected brightness in Rayleighs is shown in Figure 27, which shows a total brightness for the two multiplets (solid lines) of ~870 R at 0 km altitude, where the overhead S column density is ~$6 \times 10^{14} \text{ cm}^{-2}$. The brightnesses from the individual 147.3 and 147.41 nm lines (dashed lines) are roughly 10–50 times less. However, in contrast to the 147.9-nm multiplet (containing all transitions at 147.2–148.7 nm), the latter lines were only individually calculated, and did not include any cascade from more energetic terms. If there were only S
emissions from altitudes above ~350 km, the total UV multiplet brightness would be ~600 R, similar to the brightness as observed, from a column density of ~7 × 10¹³ cm⁻². Since optical depth = 1 occurs at N(S) ~2.5 × 10¹³ cm⁻², we do not expect to receive much radiation from below 500 km altitude.

The underlying assumption in our calculation is that when a radially outward directed photon is absorbed and reradiated, it is scattered into no preferred direction. In the lower (few 100 km) atmosphere there are most likely sufficient absorbing molecules to ensure scattered photons do not escape the atmosphere. Hence the expected ~600 R from the 147.9-nm multiplet in an atmosphere with a total S column abundance ≥6 × 10¹⁴ cm⁻² is in excellent agreement with the observed 1-µm and 0.7725-µm brightnesses. Since the various observations were taken over a >20 year period, it appears that Io's atmosphere and the plasma torus appear to be quite stable over decades-long timescales.

## 6. Non-Detections of O₂, S₂ and S₂O

**O₂:** Although photodissociation of SO₂ would also produce some O₂, the column-integrated production rate of O₂ is expected to be much less (~2 orders of magnitude) than that of SO and O, since the reaction of S + O₂ → SO + O proceeds at 2.3 × 10⁻¹² cm³ s⁻¹ so that large S and O₂ concentrations cannot coexist (Moses et al., 2002). As shown in Figure 2c, we indeed do not detect O₂ at its expected emission band near 1.27 µm. In disk-averaged spectra, the rms is ~1 × 10¹⁰ photons s⁻¹ cm⁻² µm⁻¹ sr⁻¹. Using the core of the SO line as a proxy for the core of the O₂ band, we find a 3-sigma upper limit to the brightness (4πI) of O₂ of ~1kR; that is, that is the intensity when the peak intensity of the line exceeds 3 × the rms value. With the transition probability of 1.47 × 10⁻⁴ s⁻¹ (Mlynczak & Nesbitt, 1995), this translates into an upper limit to the column density N(O₂) in the a¹Δ state of 7 × 10¹² cm⁻². Using Equation 5, with ΔE = 0.98 eV (Crisp et al., 1996), we find a column density of O₂ molecules of 8.4 × 10¹⁵ cm⁻² at 1,700 K, which is much too high an upper limit to confirm the above prediction.

**S₂O:** As mentioned in Section 4.2, in eclipse we expect the atmosphere to contain disulfur monoxide. Müller et al. (1999) showed the vibronically resolved emission spectrum of S₂O, which caught our attention since it includes a transition at 1.692 µm. However, none of the other (roughly a dozen) transitions in our wavelength range (1.60–1.86 µm) are visible in our spectra, so we do not think the 1.69 feature can be caused by S₂O.

**S₂:** Dissociation of S₂, which is isoelectronic with O₂, yields S(¹D) + S(³P₂) at 5.56 eV and is the onset of the Schumann-Runge continuum (Sun et al., 2019). Like O₂ and SO, S₂ has a forbidden a¹Δg state with the (0-0) band at 2.28 µm, 0.54 eV above the ground term X³Σ⁻g (Setzer et al., 2003). As shown in the spectra in Figure 6, we did not detect S₂ at 2.28 µm, nor above the volcanoes Kanehekili Fluctus or Loki Patera; for this band, as well as O₂, there also is no evidence of emissions on the image maps (nor after averaging image frames over the emission band to improve the signal-to-noise). As mentioned for O₂, in disk-averaged spectra, the rms is ~1 × 10¹⁰ photons s⁻¹ cm⁻² µm⁻¹ sr⁻¹. As for O₂, using the core of the SO line as a proxy for the S₂ band, we find a 3-sigma upper limit to the brightness (4πI) of the S₂ band of ~1kR. With the transition probability of 0.00286 s⁻¹ (Klotz et al., 1984), (but note that their SO transition probability was incorrect), the S₂ column density N(S₂)[a¹Δg] is ~3.5 × 10¹¹ cm⁻². Using Equation 5 for S₂, gives us an upper limit to the total S₂ column abundance of ~2 × 10¹³ cm⁻² at 1,700 K, or 1.5 × 10¹⁷ cm⁻² at 500 K.

S₂ has only been detected once on Io, in HST observations of Pele's plume against Jupiter, at a column abundance N(S₂) ~ (1 ± 0.2) × 10¹⁶ cm⁻², or a S₂/SO₂ ratio of ~0.14 (Spencer et al., 2000). Our upper limit, calculated at the atmospheric temperature of 1,700 K, is more than a factor of 100 below the Spencer et al. (2000) detected abundance of S₂. Pele falls just behind the limb of Io during its eclipse phase, and this suggests that the concentration of S₂ is highly localized to specific active plumes, as opposed to a globally distributed component of Io's atmosphere. This is supported by the red, S₃ and S₄ sulfur deposits observed at a number of sites (including Pele) where new lava is being erupted onto or exposed at Io's surface.

As a crude estimate for the required S₂ dissociation rate by electron impact one can estimate the relevant cross section from Garrett et al. (1985) as a function of incident electron velocity and when multiplied by the electron velocity one obtains a rate ~2 × 10⁻⁹ cm³ s⁻¹ for Tₑ ~ 4 – 5 eV. With a torus electron density nₑ = 2,500 cm⁻³, the dissociation rate would be ~5 × 10⁻⁶ s⁻¹. If we assume one S(¹D) is produced per dissociation event, then to
generate an emission rate of 9.65 kR for the 1 µm lines from S<sub>2</sub> dissociation above Io's northern hemisphere, a S<sub>2</sub> column density N(S<sub>2</sub>) ~9.65 × 10<sup>9</sup>/5 × 10<sup>-6</sup> ~ 2 × 10<sup>15</sup> cm<sup>-2</sup> is required, which is roughly 2 orders of magnitude above our observed disk-averaged 3-sigma upper limit at 1,700 K.

7. Conclusions

Io was observed while in eclipse with JWST's NIRSpec/IFU on 4 and 23 August 2023 in the G140H/F100LP, G235H/F170LP and G395H/F290LP settings, covering the range 1.0–5.3 µm at a resolving power R ~ 2,700. These data are compared with our first set of data from 15 November 2022 (de Pater, Lellouch, et al., 2023), when Kanehekili Fluctus was exceptionally bright, and Loki Patera had just entered a new brightening phase. The satellite was much less active in August 2023. Our main findings can be summarized as follows:

* As in November 2022, Kanehekili Fluctus was the brightest thermal source between 1 and 3 µm on the hemisphere (centered at ~345°W and 3.3°N) of Io that was observed in August 2023. At longer wavelengths, Loki Patera was the brightest source.

* By using NIR images collected with Keck close in time to the JWST observations of Io, we determined the location of the Kanehekili hot spot (28.5 ± 0.9°W and 18.2 ± 0.3°S), which was used to navigate the JWST data.

* Use of Juno/JIRAM, JunoCam, and Keck data obtained close in time to the JWST observations allowed identification of the volcanic centers at Kanehekili Fluctus versus Guaraci Fluctus. Throughout 2020–2021 both Keck and JIRAM showed Guaraci Fluctus to be brighter than Kanehekili Fluctus, though both sites were quite faint. In 2022 a brightening was observed at Kanehekili Fluctus, culminating in an exceptionally high radiance during the November 2022 JWST observations; this intensity was still considerably lower, though, than that measured in 2010 (de Pater, Davies, McGregor, et al., 2014), but higher than that observed during the Galileo epoch (Davies et al., 2012).

* The results of fitting the JWST spectra of the Kanehekili region and Loki Patera using different models (invoking different temperature-area distributions that all require silicate lava compositions) converge on a solution that reflects the changing nature of the eruptions at these locations between 2022 and 2023. In the Kanehekili region, the JWST spectra capture the likely peak of a vigorous eruption in November 2022, replacing lava flows that have cooled significantly by August 2023.

* At Loki Patera, the changes observed in the JWST spectra, model fits, determination of integrated thermal emission, and comparison with previous observations and other data show that Loki Patera had started a new brightening phase (due to lava lake resurfacing) by 22 November 2022, but thermal emission had not yet reached its peak. By August 2023, Loki Patera was in a more quiescent state, as would be expected nine months after a resurfacing episode started, and months after the period of peak thermal emission. Despite the relatively low levels of spectral radiance from Loki Patera seen by JWST in 2022 and 2023 compared to its peak thermal emission, estimates of total thermal emission (7.5–8 TW) showed Loki Patera was a significant contributor to Io's volcanic thermal emission budget.

* The 1.707 µm forbidden SO emission line complex was detected in both years; in 2022 the emission rate was 1.5 times higher than in 2023 (Table 8). In 2022 most of the SO emission came from Kanehekili Fluctus, while in 2023 the strongest SO emissions came from both the Kanehekili region and two regions in the northern hemisphere. Typical LTE thermodynamic quenching temperatures of order 1,500–1,700 K were derived based upon the total emission rates of 1.5 × 10<sup>27</sup> and 2.3 × 10<sup>27</sup> photons s<sup>-1</sup> in 2023 and 2022, respectively. The 4πl intensity, averaged over Io's disk, was 14.5 kR in 2023 and 1.5 times higher in 2022. Assuming the SO molecules above the SO-bright northern spots in 2023 to arise from unresolved vents, the 4πl brightness is of order 200 kR; with a 1,500–1,700 K temperature, this suggests SO column densities ~10<sup>15</sup> cm<sup>-2</sup> above these vents.

* The SO spectra also clearly show the 1.69 emission bump first seen in 1999 (de Pater et al., 2002), indicative of additional non-LTE processes. Since our JWST data show no exceedingly bright thermal hot spots in the northern hemisphere, this may suggest stealth volcanism as a primary source in the north. We note, though, that the absence of thermal hot spots does not always imply an absence of volcanic plumes that might be seen at visible wavelengths.

* We detected and mapped, for the first time, the forbidden [SI] 1.082 and 1.1308 µm emissions. The emissions are confined to the northern hemisphere on Io's disk, and are very smooth, not patchy as the SO emissions. Average intensities in the northern hemisphere are 7.4 and 2.25 kR at [SI] 1.082 and 1.1308 µm, respectively (i.e., 9.65 kR for the combined lines), which translates into a sulfur column density N(S[¹D])
= 3.55 × 10¹¹ cm⁻². Globally averaged numbers are a factor of 1.7 lower, that is, translate into N(S[¹D]) ~2 × 10¹¹ cm⁻².

* The observed 1.082 and 1.1308 µm brightnesses can be attained with a volcanic thermodynamic temperature of ~1,500–1,700 K and N(S) ~2 – 5 × 10¹⁴ cm⁻², ~1% of the SO₂ abundance, which is roughly a factor of 10 higher than predicted based upon equilibrium chemistry of volcanic gases above Pele (Schaefer & Fegley, 2005). This suggests that these atoms must be dissociated products in Io's atmosphere. In addition, the morphology of the emissions agrees best with that expected for electron impact excitation of atomic S as the primary excitation source. The total observed brightness in the northern hemisphere (9.65 kR) can be achieved by torus electron collisional excitation when Tₑ = 4 eV, nₑ = 2500 cm⁻³ with N(S) = 8 × 10¹⁴ cm⁻² using the effective collisional strengths computed by Tayal (2004) based on the explicit fine structure of the S(³P) term to the S (¹D and ¹S) terms. A globally averaged column density would be 4.6 × 10¹⁴ cm⁻² to provide the globally averaged total 1-µm emission of 5.6 kR.

* The electron impact calculations based on the effective collisional strengths of Tayal (2004) to explain the JWST-IR [SI] 1-µm lines also predict a brightness of ~600 R of the 0.7725 µm line, which agrees extremely well with the 0.7725-µm line recently reported by Milby et al. (2025) based upon Keck/HIRES observations (~613 R, taken two days after our JWST data, assuming the emissions are localized as in the JWST 1-µm images), and with the UV 147-nm multiplets that were observed with HST-STIS in 2000 (~590 R over the southern hemisphere). Hence, collisional excitation by electron impact from the Io plasma torus at an electron density of ~2,500 cm⁻³ and temperature of 4 eV, with an atmospheric sulfur abundance of 8 × 10¹⁴ cm⁻² produces essentially all observed [SI] emissions (1-µm, 0.7725-µm, and 147-nm). The 2023 JWST 1-µm and 2000 HST/STIS 147-nm emissions were observed over the northern and southern hemisphere, respectively, as expected.

* Based upon calculations of electron impact excitation on sulfur, and the requirement that the population levels ³P₁₀ need to be sufficiently populated to excite enough S atoms to the ¹D level, the atmospheric temperature should be roughly 1,700 K.

* Disk-averaged 3-σ upper limits were obtained for O₂, S₂O, and S₂ of ~1 kR. For S₂ this translates into a column abundance N(S₂) ≲2 × 10¹³ cm⁻² at 1,700 K, which is 2 orders of magnitude below both the column abundance observed above Pele (Spencer et al., 2000), and the column abundance needed to generate 9.65 kR for the 1 µm [SI] lines from S₂ dissociation above Io's northern hemisphere. At lower temperatures, the N(S₂) upper limit is higher.

The fact that we can explain the sulfur emissions in the IR (~1-µm) with JWST-NIRSpec, visible (0.7725-µm) with Keck-HIRES taken between 2022 and 2024, and in the UV (147.9-nm multiplets) with HST-STIS taken 20 years earlier, with the same plasma torus electron density (~2,500 cm⁻³), temperature (4 eV), and sulfur column abundance (8 × 10¹⁴ cm⁻²) implies that the system (plasma torus and Io's atmosphere) is quite stable over decades-long timescales. To further investigate the stability of this system and the effect of volcanic eruptions on the emissions, observations of the 1-µm and 0.7725-µm lines plus the 147.9-nm multiplets need to be carried out at the same time to ensure the same viewing geometry and volcanic activity.

## Appendix A: Keck Observations

The satellite Io was observed with the NIRC2 camera coupled to the Adaptive Optics (AO) system on the Keck II telescope at the W.M. Keck Observatory in Hawaii (Wizinowich et al., 2000) as summarized in Tables 2 and 6. NIRC2 is a 1,024 × 1,024 Aladdin-3 InSb array, which was used in its highest angular resolution mode, that is, the NARROW camera, with a plate scale (after 2015) at 9.971 ± 0.004 mas per pixel (Service et al., 2016), which translates roughly into 40 km/pixel at the center of Io's disk when the satellite is at a geocentric distance of 5.5 AU, down to 32 km/pixel for a geocentric distance of 4.0 AU. Io itself was used for wavefront sensing. All images were processed using standard near-infrared data reduction techniques (flat-fielded, sky-subtracted, with bad pixels replaced by the median of surrounding pixels). The geometric distortion in the Keck images was corrected using the solution provided by Service et al. (2016). The individual images were aligned and co-added to increase signal to noise. Photometric calibration was performed on the standard stars HD201941, HD19600, and HD1160 on 15 August 2022, 7, 12 September, and 22 June 2023 (Elias et al., 1982; Leggett et al., 2003); on non-photometric nights or nights where no stars were observed, the photometry was bootstrapped from well-calibrated images while properly accounting for the variations in Io's heliocentric and geocentric distance at
the time of the observations (as in, e.g., de Pater, Davies, McGregor, et al., 2014). Two observations were taken during our Twilight Zone Program (de Pater, de Kleer, Molter, et al., 2017): 17 April 2021, and 11 August 2023.

## Appendix B: UV Calculations

To compute the brightnesses of the allowed $S(^3D^0-^3P_2$, 147.41 nm) and semi-forbidden $S(^5D^0-^3P_2$, 147.3 nm) lines in the thin Io atmosphere, we assume that they have Doppler line profiles defined in terms of frequency $\nu_0$ with $\delta\nu = \nu - \nu_0$ and given by

$$ \phi(\delta\nu) = \left(\frac{1}{\delta\nu_D\sqrt{\pi}}\right) \exp\left[-\left(\frac{\delta\nu}{\delta\nu_D}\right)^2\right], \qquad (B1) $$

with line center cross-section $\sigma(\nu_0) = \mathcal{L}\phi(\nu_0) = \sigma_0$, and $\mathcal{L}$ is the line intensity, and the line cross-section is a function of frequency as defined by $\sigma(\delta\nu) = \mathcal{L}\phi(\delta\nu)$. $\delta\nu_D$ is the Doppler width of the line for a temperature $T = 1,700$ K, and $\nu_0$ as given below,

$$ \delta\nu_D = \left(\frac{\nu_0}{c}\right) \left(\frac{2kT}{m}\right)^{0.5} = 6.3775 \times 10^9 \text{ Hz.} \qquad (B2) $$

The brightnesses of these UV lines are then given by the integral of the product of the excitation rate ($\Gamma$) times the sulfur density ($n_S$) times the electron density ($n_e$) with the added factor of the Doppler line shape times the transmission of Doppler lines through the overhead column of sulfur atoms. This is carried out over an effective line width of 8 Doppler widths. In equation form it is:

$$ 4\pi I(T_e, z) = 10^{-6} \int_{-4, \delta\nu_D}^{4, \delta\nu_D} \int_{z}^{\infty} \Gamma(\text{SI})[147 \times X(T_e)] \cdot n_e \cdot n_S(z) \cdot \phi(\delta\nu) \cdot \exp \left[ - \int_{z}^{\infty} (\mathcal{L} \cdot \phi(\delta\nu) \cdot n_S(z_l) \cdot dz_l) \cdot dz \right] \cdot d\delta\nu, \qquad (B3) $$

where 147.X stands for the 147.41 and 147.3 nm lines. For the allowed $S(^3D^0-^3P_2$, 147.41 nm) line, $\nu_0 = 2.0351 \times 10^{15}$ s$^{-1}$, $\sigma_0 = 2.103 \times 10^{-13}$ cm$^2$, $\mathcal{L} = 2.3774 \times 10^{-3}$ cm$^{-2}$. For the semi-forbidden $S(^5D^0-^3P_2$, 147.3 nm) line, $\nu_0 = 2.0352 \times 10^{15}$ s$^{-1}$, $\sigma_0 = 3.9963 \times 10^{-14}$ cm$^2$, and $\mathcal{L} = 4.5173 \times 10^{-4}$ cm$^{-2}$.

## Appendix C: Tables of Collisional Strengths, Excitation Rates, and the Effect of Atmospheric Temperature

To understand the effect of temperature on S emissions due to electron impact excitation on the ¹S and ¹D states, one needs to examine the boost that the Tayal (2004) method gives from excitations mostly from the ³P₁ state which is 0.049 eV above the ³P₂ state and the ³P₀ state which is 0.071 eV above the ³P₂ state (Figure 22). Referring to Tables C1 and C2 for the collision excitation rates to ¹D from [³P₂, ³P₁, ³P₀] at the critical $T_e = 4.32$ eV (4.9 × 10$^{-10}$, 6.53 × 10$^{-09}$, 3.3 × 10$^{-08}$ cm$^3$ s$^{-1}$) and to ¹S from [³P₂, ³P₁, ³P₀] also at $T_e = 4.32$ eV (5.3 × 10$^{-11}$, 5.64 × 10$^{-10}$, 2.61 × 10$^{-09}$ cm$^3$ s$^{-1}$), one can see the critical role that elevated atmospheric temperatures play in populating ³P₁ and ³P₀ from which the excitation rates far exceed the ³P₂ rate in spite of its statistical weight of 5 relative to the respective asymptotic values of 3 and 1 for ³P₁ and ³P₀ states. The older Zatsarinny and Tayal (2002) method does not populate the ³P levels according to thermodynamic equilibrium considerations at the relevant atmospheric temperature, which results in the asymptotic statistical weights of 5, 3, 1. When using this older method to match the observed 1-µm lines, the S column density would be 3.6 × 10$^{14}$ cm$^{-2}$ ($n_e = 2,500$ cm$^{-3}$), and the 0.7725 µm brightness would be ~800 R, that is, ~30% higher than observed and predicted with the more recent Tayal (2004) method.
<caption>Table C1</caption>

Effective Collisional Strengths (γ) and Excitation Rates for Excitation From S($^3P_{1,2,0}$) Multiplets to S($^1D$) as a Function of Electron Temperature in eV

<table><caption>Table C2</caption><thead><tr><th>T<sub>e</sub><br>(K)</th><th>T<sub>e</sub><br>(eV)</th><th>γ(1130)</th><th>E/1.096<br><sup>3</sup>P<sub>1</sub>-<sup>1</sup>D</th><th>Rate<br>(cm<sup>3</sup>/s)</th><th>γ(1082)</th><th>E/1.146<br><sup>3</sup>P<sub>2</sub>-<sup>1</sup>D</th><th>Rate<br>(cm<sup>3</sup>/s)</th><th>γ(1154)</th><th>E/1.074<br><sup>3</sup>P<sub>0</sub>-<sup>1</sup>D</th><th>Rate<br>(cm<sup>3</sup>/s)</th></tr></thead><tbody><tr><td>1,995</td><td>0.172</td><td>0.035</td><td>0.157</td><td>3.85E-12</td><td>0</td><td>0.1501</td><td>0</td><td>0.058</td><td>0.1602</td><td>2.18E-11</td></tr><tr><td>2,512</td><td>0.217</td><td>0.044</td><td>0.198</td><td>1.60E-11</td><td>0</td><td>0.1890</td><td>0</td><td>0.073</td><td>0.2016</td><td>8.82E-11</td></tr><tr><td>3,162</td><td>0.273</td><td>0.053</td><td>0.249</td><td>4.87E-11</td><td>0.002</td><td>0.2379</td><td>9.17E-13</td><td>0.088</td><td>0.2538</td><td>2.63E-10</td></tr><tr><td>3,981</td><td>0.343</td><td>0.071</td><td>0.313</td><td>1.33E-10</td><td>0.002</td><td>0.2995</td><td>1.94E-12</td><td>0.118</td><td>0.3195</td><td>7.06E-10</td></tr><tr><td>5,012</td><td>0.432</td><td>0.088</td><td>0.394</td><td>2.83E-10</td><td>0.003</td><td>0.3770</td><td>5.16E-12</td><td>0.146</td><td>0.4022</td><td>1.48E-09</td></tr><tr><td>7,943</td><td>0.685</td><td>0.135</td><td>0.625</td><td>8.80E-10</td><td>0.004</td><td>0.5975</td><td>1.45E-11</td><td>0.225</td><td>0.6376</td><td>4.54E-09</td></tr><tr><td>10,000</td><td>0.862</td><td>0.164</td><td>0.787</td><td>1.32E-09</td><td>0.006</td><td>0.7522</td><td>2.74E-11</td><td>0.273</td><td>0.8027</td><td>6.78E-09</td></tr><tr><td>19,953</td><td>1.720</td><td>0.299</td><td>1.569</td><td>3.22E-09</td><td>0.013</td><td>1.5009</td><td>8.16E-11</td><td>0.501</td><td>1.6015</td><td>1.64E-08</td></tr><tr><td>31,623</td><td>2.726</td><td>0.431</td><td>2.487</td><td>4.67E-09</td><td>0.03</td><td>2.3788</td><td>1.91E-10</td><td>0.724</td><td>2.5383</td><td>2.37E-08</td></tr><tr><td>34,811</td><td>3.432</td><td>0.551</td><td>3.131</td><td>5.77E-09</td><td>0.055</td><td>2.9947</td><td>3.41E-10</td><td>0.929</td><td>3.1955</td><td>2.94E-08</td></tr><tr><td>50,119</td><td>4.321</td><td>0.655</td><td>3.942</td><td>6.53E-09</td><td>0.083</td><td>3.7701</td><td>4.90E-10</td><td>1.106</td><td>4.0229</td><td>3.33E-08</td></tr><tr><td>63,096</td><td>5.439</td><td>0.742</td><td>4.963</td><td>6.95E-09</td><td>0.11</td><td>4.7463</td><td>6.12E-10</td><td>1.254</td><td>5.0645</td><td>3.54E-08</td></tr><tr><td>100,000</td><td>8.621</td><td>0.952</td><td>7.866</td><td>7.63E-09</td><td>0.192</td><td>7.5224</td><td>9.18E-10</td><td>1.609</td><td>8.0267</td><td>3.88E-08</td></tr><tr><td>199,526</td><td>17.201</td><td>1.033</td><td>15.694</td><td>6.24E-09</td><td>0.257</td><td>15.0092</td><td>9.29E-10</td><td>1.744</td><td>16.0154</td><td>3.17E-08</td></tr><tr><td>316,228</td><td>27.261</td><td>0.957</td><td>24.873</td><td>4.70E-09</td><td>0.254</td><td>23.7880</td><td>7.48E-10</td><td>1.612</td><td>25.3827</td><td>2.38E-08</td></tr><tr><td>398,107</td><td>34.320</td><td>0.865</td><td>31.313</td><td>3.82E-09</td><td>0.237</td><td>29.9473</td><td>6.27E-10</td><td>1.454</td><td>31.9549</td><td>1.93E-08</td></tr></tbody></table>

Effective Collisional Strengths (γ) and Excitation Rates for Excitation From S($^3P_{1,2,0}$) Multiplets to S($^1S$) as a Function of Electron Temperature in eV

<table><thead><tr><th>T<sub>e</sub><br>(K)</th><th>T<sub>e</sub><br>(eV)</th><th>γ(4589)</th><th>E/2.70<br><sup>3</sup>P<sub>1</sub>-<sup>1</sup>S</th><th>Rate<br>(cm<sup>3</sup>/s)</th><th>γ(4507)</th><th>E/2.75<br><sup>3</sup>P<sub>2</sub>-<sup>1</sup>S</th><th>Rate<br>(cm<sup>3</sup>/s)</th><th>γ(4628)</th><th>E/2.679<br><sup>3</sup>P<sub>0</sub>-<sup>1</sup>S</th><th>Rate<br>(cm<sup>3</sup>/s)</th></tr></thead><tbody><tr><td>1,995</td><td>0.172</td><td>0.0040</td><td>0.0637</td><td>3.92E-17</td><td>0</td><td>0.0625</td><td>0</td><td>0.007</td><td>0.0642</td><td>2.33E-16</td></tr><tr><td>2,512</td><td>0.217</td><td>0.0050</td><td>0.0802</td><td>1.10E-15</td><td>0</td><td>0.0787</td><td>0</td><td>0.009</td><td>0.0808</td><td>6.56E-15</td></tr><tr><td>3,162</td><td>0.273</td><td>0.0060</td><td>0.1010</td><td>1.53E-14</td><td>0</td><td>0.0991</td><td>0</td><td>0.01</td><td>0.1018</td><td>8.28E-14</td></tr><tr><td>3,981</td><td>0.343</td><td>0.0080</td><td>0.1271</td><td>1.40E-13</td><td>0</td><td>0.1248</td><td>0</td><td>0.014</td><td>0.1281</td><td>7.80E-13</td></tr><tr><td>5,012</td><td>0.432</td><td>0.0100</td><td>0.1600</td><td>7.84E-13</td><td>0</td><td>0.1571</td><td>0</td><td>0.017</td><td>0.1613</td><td>4.20E-12</td></tr><tr><td>7,943</td><td>0.685</td><td>0.0150</td><td>0.2536</td><td>9.39E-12</td><td>0.001</td><td>0.2490</td><td>3.49E-13</td><td>0.026</td><td>0.2556</td><td>5.03E-11</td></tr><tr><td>10,000</td><td>0.862</td><td>0.0190</td><td>0.3193</td><td>2.38E-11</td><td>0.001</td><td>0.3135</td><td>7.10E-13</td><td>0.031</td><td>0.3218</td><td>1.20E-10</td></tr><tr><td>19,953</td><td>1.720</td><td>0.0360</td><td>0.6371</td><td>1.53E-10</td><td>0.002</td><td>0.6255</td><td>4.94E-12</td><td>0.06</td><td>0.6421</td><td>7.72E-10</td></tr><tr><td>31,623</td><td>2.726</td><td>0.0530</td><td>1.0097</td><td>3.18E-10</td><td>0.005</td><td>0.9913</td><td>1.77E-11</td><td>0.089</td><td>1.0176</td><td>1.62E-09</td></tr><tr><td>39,811</td><td>3.432</td><td>0.0680</td><td>1.2711</td><td>4.46E-10</td><td>0.009</td><td>1.2480</td><td>3.49E-11</td><td>0.114</td><td>1.2811</td><td>2.26E-09</td></tr><tr><td>50,119</td><td>4.321</td><td>0.0820</td><td>1.6002</td><td>5.64E-10</td><td>0.013</td><td>1.5711</td><td>5.30E-11</td><td>0.126</td><td>1.6128</td><td>2.61E-09</td></tr><tr><td>63,096</td><td>5.439</td><td>0.0920</td><td>2.0146</td><td>6.41E-10</td><td>0.017</td><td>1.9779</td><td>7.04E-11</td><td>0.154</td><td>2.0303</td><td>3.23E-09</td></tr><tr><td>100,000</td><td>8.621</td><td>0.1170</td><td>3.1928</td><td>7.78E-10</td><td>0.027</td><td>3.1348</td><td>1.07E-10</td><td>0.195</td><td>3.2179</td><td>3.90E-09</td></tr><tr><td>199,526</td><td>17.201</td><td>0.1240</td><td>6.3706</td><td>8.82E-10</td><td>0.033</td><td>6.2547</td><td>1.09E-10</td><td>0.207</td><td>6.4205</td><td>3.42E-09</td></tr><tr><td>316,228</td><td>27.261</td><td>0.1120</td><td>10.0967</td><td>5.19E-10</td><td>0.032</td><td>9.9131</td><td>8.88E-11</td><td>0.188</td><td>10.1758</td><td>2.61E-09</td></tr><tr><td>398,107</td><td>34.320</td><td>0.1000</td><td>12.7110</td><td>4.21E-10</td><td>0.029</td><td>12.4798</td><td>7.32E-11</td><td>0.167</td><td>12.8106</td><td>2.11E-09</td></tr></tbody></table>
## Acknowledgments

This work is based on observations made with the NASA/ESA/CSA James Webb Space Telescope. The data were obtained from the Mikulski Archive for Space Telescopes at the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-03127 for JWST. These observations are associated with program #1373 (Observations #35 and #36), which is led by co-PIs Imke de Pater and Thierry Fouchet. Some of the data presented in this paper were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation. The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Maunakea has always had within the indigenous Hawaiian community. Without their generous hospitality, we would not have been able to obtain the observations presented here. We thank Bryan Holler for his help/advice on the data reduction. We further thank Edward Molter and the Keck Support Astronomers Carlos Alvarez and Randy Campbell for their help with the Twilight Zone Project. IdP and MHW are in part supported by the Space Telescope Science Institute grant nr. JWST-ERS-01373. TF and EL acknowledge support from ANR Grant ANR-21-CE49-0020-01. LNF was supported by STFC Consolidated Grant reference ST/W00089X/1. KdK acknowledges support from NASA through Grant HST-GO-15425.002-A from the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5-26555. CS acknowledges support from NASA programs 80NSSC22K0954, 80NSSC21K1138 and HST-GO-17470. J.S. and S.S. received funding from the European Research Council (ERC) under the European Union's Horizon 2020 research and innovation programme (grant agreement No. 884711). Part of this work was performed at the Jet Propulsion Laboratory, California Institute of Technology, under Government contract. AGD, JEP, DAW, and DMN are supported by award 80NM0018F0612 from the NASA New Frontiers Data Analysis Program. For the purpose of open access, the author has applied a Creative Commons Attribution (CC BY) license to the Author Accepted Manuscript version arising from this submission.

## Conflict of Interest

The authors declare no conflicts of interest relevant to this study.

## Data Availability Statement

The JWST data were obtained from the Mikulski Archive for Space Telescopes at the Space Telescope Science Institute, Observations #35 and #36 of ERS program #1373. The specific observations analyzed can be accessed via de Pater (2024). The Keck data were obtained from the Keck Archive (https://koa.ipac.caltech.edu). Juno/JIRAM data are published by Perry et al. (2025).

The data were reduced with the standard JWST pipeline (https://jwst-docs.stsci.edu/jwst-science-calibration-pipeline-overview) and a simple IDL command to correct the data for the 1/f noise, as described in Section 2 (the 1/f noise correction can now be done using the latest JWST pipeline). The resulting fits files were read into IDL, and aligned by using the Fast Fourier transform (FFT) cross-correlation algorithm. These data cubes were inspected and numerous cosmic rays and bad pixels were removed (mostly by hand), as explained in the text. To make spectra the authors used IDL routine apphot.pro (Perrin, 2013), after expanding the images by a factor of 5, and interpolating to the smaller-sized pixels.

The final cleaned spectral data cubes are uploaded to Zenodo. de Pater (2025a) contains the cleaned data cubes from the November 2022 dataset. de Pater (2025b) contains the data sets from 4 August 2023. de Pater (2025c) contains the data sets from 23 August 2023. Thermal emission, SO and [SI] spectra as displayed in the various figures can be retrieved from de Pater (2025d).

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