Detection of the infrared aurora at Uranus with Keck-NIRSPEC

Near infrared (NIR) wavelength observations of Uranus have been unable to locate any infrared aurorae, despite many attempts to do so since the 1990s. While at Jupiter and Saturn, NIR investigations have redefined our understanding of magnetosphere ionosphere thermosphere coupling, the lack of NIR auroral detection at Uranus means that we have lacked a window through which to study these processes at Uranus. Here we present NIR Uranian observations with the Keck II telescope taken on the 5 September 2006 and detect enhanced $\text{H}_{\text{3}}^{\text{+}}$ emissions. Analysing temperatures and column densities, we identify an 88\% increase in localized $\text{H}_{\text{3}}^{\text{+}}$ column density, with no significant temperature increases, consistent with auroral activity generating increased ionization. By comparing these structures against the $\text{Q}_{\text{3}}^{\text{mp}}$ magnetic field model and the Voyager 2 ultraviolet observations, we suggest that these regions make up sections of the northern aurora.

auroral, the spectra were analysed for temperature, column density and total emissions to identify whether enhancements were thermally driven or created by an ion population increase.
Uranus observations were taken with the Keck II telescope on the 5th September 2006, from 07:26 to 13:24 UT, close to the planet's equinox in 2007, using the NIRSPEC (Near-infrared Spectrograph) instrument [27] with a KL atmospheric window filter.A 0.288 × 24 arcsec slit was aligned with the plant's rotational axis (shown in Fig. 1a).Spectra were gathered between 3.5 µm and 4.1 µm where the fundamental Q-branch of H + 3 emissions lies (shown in Fig. 1b; raw image in Extended Data Figure 1).This triatomic hydrogen ion is a major constituent of Uranus's ionosphere and planets whose upper atmosphere is dominated by molecular and ionic hydrogen [28].A total of 218 spectra were taken over an 6 h period with an integration time of 30 s.These were co-added into 13 datasets to enhance the signal-to-noise ratio (total integration time per set was 27 min).To increase the signal-to-noise ratio further, spatial pixels along the slit were grouped every 0.32 arcsecs (full details in the Methods).The exact longitude of Uranus has been completely lost; therefore, an arbitrary longitude has been selected for these results.Astronomical seeing on the night averaged at 0.44 arcsec, which is equivalent to a blur of 14°latitude and 12°longitude.During the observation, Uranus rotated by ∼180°and hence our final mapping spans an area up to ∼180°longitude.Unfortunately, a lapse in guiding between 10:52 UT and 11:31 UT resulted in the loss of 2 longitudinal data bins, leaving a gap in the middle of our scans.Finally, results presented here are not corrected for line of sight (LOS) (for example, see [17]) and hence we expect infrared emissions to be enhanced near the planet's limb.At Jupiter and Saturn, auroral emissions are LOS enhanced; however, Uranus's solar extreme ultraviolet (EUV)-generated ionosphere is darker at the limbs [7], and, so, without a detailed understanding of the ionospheric brightening source, it is not possible to correct.However, as much of the enhanced emissions are away from the limbs, we expect minimal change in the location of emissions peaks after corrections.To calculate the H + 3 intensities, temperatures, column densities and total H + 3 emission for the upper atmosphere of Uranus, this study focuses on five quasi-thermalized ro-vibrational emission lines of H + 3 , Q(1,0 -), Q(2,0 -), Q(3,0 -), Q(3,1 -) and Q(3,2 -); these physical parameters were calculated from a full spectrum best fit, as described in Methods.The final fitted spectra provide intensity values that are then mapped across Uranus as is shown with the Q(1,0 -) emission line (with the highest signal-to-noise ratio; Fig. 2a), H + 3 total emission (Fig. 2b), temperatures (Fig. 2c) and H + 3 column density (Fig. 2d).The respective error maps are shown in Extended Data Figure 2.
In Fig. 2a, b, the H + 3 emission intensity varies with local time.To confirm the source of these enhancements, we define three regions of interest that are algorithmically distinct: the 'enhanced' region where the emissions are brighter than the mean plus one standard deviation (shown in solid black lines but not shaded); the 'dim' region where emissions are below the mean emission (shaded with dots); and the 'intermediate' region where emissions are brighter than the mean, but within a standard deviation of that mean (shaded by diagonal lines).The means and standard deviations for Fig. 2 presented in Table 1 are the result of subtracting each pixel by its uncertainty (seen in Extended Data Figure 2).The resulting datasets are hence minimized, meaning pixels in the enhanced region are statistically significant.
In Fig. 2a, the enhanced regions show intriguing structures, the first, which is smaller, between 26°S and 59°S and from 18°to 28°longitude (E1).The second area extends between 15°N and 75°N from 100°and 143°longitude with two smaller emission spots between 10°N and 0°and between 10°S and 20°S over a 108°to 117°longitude range, which we refer to as E2.Table 1 summarizes the mean values of Q(1,0 -) intensities, along with mean values for temperature, column density and total Q(1,0 -) emission.Comparing the dim region's mean Q(1,0 -) intensity (0.472 ± 0.086 µWm -2 sr -1 ) with that of the enhanced region (0.723 ± 0.010 µWm -2 sr -1 and 0.716 ± 0.009 µWm -2 sr -1 , we find a 27 % to 90 % enhancement. Figure 2b shows the total H + 3 emission, which is the combined intensity from all H + 3 emission lines in this investigation.We find the two enhanced regions average at 6.155 ± 0.681 µWm -2 sr -1 and 6.354 ± 0.616 µWm -2 sr -1 for E1 and E2, respectively, while the dim region emits at a lower average of 3.212 ± 1.235 µWm -2 sr -1 , hence an 18 %, up to 353 % increase at both E1 and E2.This large range in emission enhancement is most likely from the high uncertainty in column density, which affects the error propagation when calculating the total emission.We, however, conclude that our division of emissions of Uranus-whether the single Q(1,0 -) line or total H + 3 emission, into distinct enhancement related regions, is both robust and significant.

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Mean Q( Comparing physical parameters between enhanced regions provides an understanding of how they are enhanced.The average temperature for the dataset is 585 ± 25 K, which aligns with previous temperature observations ( [25] for 2006 at 608 ± 12 K).The enhanced regions have a mean temperature of 585 ± 14 K and 564 ± 22 K for E1 and E2, respectively, with the dim region temperature at 593 ± 24 K, shown in Fig. 2c.While the enhanced regions appear cooler, there is overlap in temperature errors, so while we cannot conclude the emission is anticorrelated with temperature, thermal processes cannot explain the intensity enhancements.
Except at the planet's limbs, EUV ionization produces a uniform column ionization rate across the whole disk; however, enhancements of column densities could be produced by enhanced particle precipitation, suggestive of auroral activity.Shown in Fig. 2d, we observe an average column-density difference of 2.133 × 10 15 m -2 at the enhanced regions (4.017 ± 0.457 × 10 15 m -2 and 5.113 ± 0.826 × 10 15 m -2 at E1 and E2, respectively) compared with the dim region (2.432 ± 0.901 × 10 15 m -2 ).These densities, on average, are higher (about two to five times higher) than reported in [25].Here a more through and complete data-reduction process was conducted over the whole night of observations rather than half the night, with densities presented in Table 1 aligning within the range of densities observed in previous investigations [25], [26].In contrast to the temperatures, the enhanced region's column density is on average 88% enhanced.Put simply, more emitters, rather than hotter emitters, is what is resulting in more emission.
There are several scenarios that could lead to a column-density enhancement at locations of increased H + 3 emissions.One possibility is if the ion is produced evenly across the planetary disk, there is some mechanism by which it is transported from the dim region into the enhanced regions.
We do not consider meridional transport from the rotational poles to be substantial for two reasons: first, Uranus is a large, rapidly rotating planet where it is difficult to overcome the Coriolis forces; second, if there are equatorwards winds, we would expect to see a H + 3 bulge evenly distributed at lower latitudes.There is nothing in our data to suggest polewards meridional transport.
Zonal winds on Uranus are generally between 0 and 250 ms -1 .A previous study [29] found electron densities between ≈ 10 9 m -3 (Voyager egress) and ≈ 10 10 / 10 11 m -3 (Voyager ingress).Taking these figures together with the dissociative recombination coefficient of ≈ 10 -13 m 3 s -1 ( [30]) suggests a maximum half lifetime τ(H + 3 ) of less than 10 4 s, and possibly as low as 100 s.Hence an individual H + 3 ion could be transported ≈ 2,000 km.This is less than the 30,000 km at the equator to get from the centre of E2; although the distance from there to the centre of E1 is approximately half that value, it is still too far.Hence, we assume that the H + 3 ions, their emissions and physical parameters are representative of locally produced features.Another potential driver for the dim region's low column densities could be 'ring rain' as seen at Saturn [31].Here, H + 3 destruction is modulated by water molecules in the planet's rings travelling along the field lines into the planet's lower latitudes.Fig. 3a combines Fig. 2a and the Q 3 model from [32], which maps Uranus's magnetic field with dip angle contours, using contour steps of 20 °dip angle (the angle made with the planet's horizontal plane by its magnetic-field lines).We expect the ring rain to affect only a narrow band of dip angles (mapping to the planet's rings, 1.6-2 R U ), where in Fig. 3a we observe the dim region over a large range of dip angles.Hence, quenching ring rain cannot explain the emissions we observe.
Two more magnetic-field models Q mp 3 and AH 5 ; [4]) have since been used at Uranus, replacing previous models with a more globally representative magnetopause image field and including UV auroral emissions from Voyager II, respectively.These models provide a strong fit to the southern aurora, but the northern aurora is poorly constrained as Voyager crossed magnetic-field lines that mapped close to the southern magnetic pole twice, once at a distance of 4.19 R U , but only once at the north, at >20 R U .In addition, the auroral morphology may have changed with solar-wind pressure or by changes in the preferred auroral acceleration region above the planet.Given this complexity, we focus solely on the Q 3 model.
As none of the previous processes can explain the NIR enhancement morphology, the most plausible explanation is that the den-4/12 Figure 3. a, Q(1,0 -) mapped intensities with the Q 3 model with contours representing the contours of 20°dip angles to the thick continuous white line, which is the magnetic-field equator.Here the ULS has been sourced and placed into our observations from the Q 3 model.The grey background colour represents the areas that were unobserved in these observations.b, Q(1,0 -) intensities (where only the enhanced region has been highlighted) mapped alongside the L shells of 2 (solid line), 3 (dashed), 5 (dotted), 10 (dot and dashed) and 20 (solid) of the Q mp 3 model.c, Q(1,0 -) intensities (where only the enhanced region has been highlighted) mapped alongside the H2 band emissions intensity map from [4] To avoid obscuring the UV emissions, the dark grey background has been removed in this panel.sity enhancements are driven by auroral production.In previous H + 3 investigations at Jupiter and Saturn [11], [12], [13], [15], [17], [18], we have observed the strongest infrared emission enhancements and column densities at the auroral regions, where particle precipitation results in significant ionization in the upper atmosphere.We find that the enhanced regions strongly suggest auroral production and so consider that we have partially mapped the northern infrared aurora.
To determine how the 2006 infrared emissions aligned with previous models, we have chosen to not add a longitude shift into our work due to the lack of known longitude (ULS) in 2006.Comparing Fig. 2a and the Q 3 model (Fig. 3a), we observe intensities between 60 °and 80 °dip angle at the same angle as where the auroral oval sits, although the approximate location of the auroral oval in the Q 3 model sits within the dim region (similar in location to Fig. 3b).Other enhanced regions with poor alignment (where the dip angle drops to 20 °) may be due to more complex morphology within the surface magnetic-field structure, or the effects of seeing (at least ≈ ±12 °) along with the low spatial resolution (≈ 0.32 arcsec).We also note relatively weak emissions between 40 °and 100 °longitude.While most pixels in this region are 'dim', we highlight that these emissions remain brighter than the limbs.It may be that while the enhancement is not significant as the enhanced region, only the edges of the map represent the EUV ionized background H + 3 density.This weaker central region could be driven by weaker auroral

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precipitation and hence further investigations are required.
In Fig. 3b, we compare the infrared emissions against the L-shell magnetic-field lines of the Q 3 model.Here the emissions observed in both E1 and E2 extend out past the optimal L5 shell, which is where the brightest UV emissions are observed by Voyager II and HST.Focusing on E1, between 30°S and 60°S and before 30°ULS longitude, we find no enhanced emissions align with the Q 3 L shells.These emissions are, however, located close to dayside O-source radio emissions [33], and are close to n-smooth radio emissions observed in [10], where the authors suggested that these emissions arose from unusual particle distribution from particle absorption by the ε ring, which may act as a driver for these infrared extended emissions.We do, however, find a portion of E2 emissions fit within L shells of 3 and 5, where weak UV emissions in Fig. 3c are located (≈ <100 R).
Figure 3c compares Fig. 2a with UV auroral emissions from Voyager II in 1986.At Jupiter, UV and infrared aurora appear at similar latitudes [34], [35]; however, at the auroral oval, [36] found UV and infrared auroral features' brightness can vary independently of each other and hence are not co-located.We should then not expect the brightest NIR emissions to be co-located with the brightest UV emissions at Uranus.Further differences between infrared and UV emissions can also be explained by the ≈ 15 min lifetime of H + 3 (at Jupiter) smoothing out short-term (1-2 min) variability in UV emissions [37].A similar effect may also occur at Uranus.
The enhanced H + 3 emissions are broadly spaced in latitude compared with the brightest UV emission, where the strongest UV emissions occur north of E1 and only weaker UV emissions appear close to or at E2.This spreading of infrared emissions suggests that H + 3 emissions occur more equatorwards (magnetic-field equator) than the UV emissions and appear anticorrelated in terms of longitude.Differences in emission region may result from changes in the auroral drivers, changes in the solar wind (as observed at Jupiter and Saturn [38]); or short-term variability associated with the local time.These might be via changing precipitation flux or precipitation energy.Equally, contrasts in apparent magnetic mapping of the two aurorae could originate from poor alignment of our arbitrary longitude.It is difficult to draw too many conclusions without re-discovering the rotational phase of Uranus.This likely detection of H + H + 3 total emission calculations Using the calculations of [46], the total emission can be calculated by the product of the number of ions by the temperaturedependent total emission per molecule (Emol) while assuming local thermal equilibrium.This requires both the column density and temperature over two or more emission lines), where temperature is used to calculate Emol.It should be highlighted that due to temperatures staying between 500 K ≤ T ≥ 900 K, suitable coefficient values were selected to calculate Emol.

Figure 2 .
Figure 2. a, Measured H + 3 Q(1,0 -) intensity mapped across the upper atmosphere of Uranus against Uranian latitude and arbitrary longitude.b, Total H + 3 emission calculated from the temperature and column density (explained in detailed in Methods).c, Estimated temperatures of the H + 3 emissions from all five Q-branch lines.d, Estimated column densities of H + 3 emissions from all five Q-branch lines.The latitude is planetocentric whereas the longitude is arbitrary due to the loss of the ULS since Voyager II.The solid black lines mark out the boundaries of E1 (left) and E2 (right).Within the boundaries, the enhanced regions are unshaded, the dim regions are shaded with dots and the intermediate regions are shaded with diagonal lines.

Table 1 .
Means and standard deviations of the Q(1,0 -) intensities, H + 3 ion temperatures, ion column densities and total emission for the enhanced and dim regions