Atmospheric Science Questions for a Uranian Probe

The mechanisms that formed the Solar System, and the Ice Giants in particular, are active areas of study (e.g., Atreya et al (2020)). Competing theories of planetary formation mechanisms fall into four main categories: the currently more generally-favored model of core accretion (Pollack et al, 1996; Hubickyj et al, 2005), the theory of disk instability (Boss, 1997; Boss et al, 2002), the photoevaporation model (Guillot and Hueso, 2006), and the CO snowline model (Ali-Dib et al, 2014). Since the composition and temperature of the protosolar nebula was highly radially and temporally dependent, the abundances of certain diagnostic species and their isotopes relative to solar, cometary, or other planetary abundances can greatly aid in constraining the conditions of Uranus’s formation. Especially important are the abundances of noble gases, which can inform the nature of the formation process and whether it occurred in a step-wise fashion or more instantaneously and can only be measured in situ.

The process of core accretion first necessitates the condensation of solid and/or icy planetesimals and the subsequent, rapid infall of large amounts of gas onto that solid core once it becomes massive enough. The possibility of this mechanism was originally limited by the length of its predicted time frame, which was at or longer than the lifetime of the protosolar nebula. However, more recent models have found that this process could take place on timescales short enough to render it possible and even likely (Rice, 2022). Depending on the way in which volatiles were incorporated into the solids that eventually feed condensing planitesimals, different compositional signatures will be reflected in the planet’s modern bulk abundances (Mousis et al, 2020). Under the assumption that there is no process leading to differences in fractionation between volatiles, the theory of core accretion via amorphous ice predicts homogeneous enrichments of O, C, N, S, Ar, Kr and Xe since their trapping efficiencies within amorphous ice are all very similar (Owen et al, 1999; Bar-Nun et al, 2007). In the case of core accretion by way of clathrates, a strong variation of trapping efficiencies between species is implied (Mousis et al, 2010), resulting in elemental abundances that vary as a function of both trapping temperature and dependence upon the availability of crystalline water, which is necessary to achieve either full or partial clathration (Gautier et al, 2001; Mousis et al, 2014, 2012).

The disk instability model of planet formation predicts that instead of a more gradual condensation of solids and subsequent rapid collapse of gas, a gravitational instability within the protosolar disk causes a self-gravitating collapse of material into a planet much more quickly, on the order of $\sim$1000 years (Boss et al, 2002). When applied to our Solar System, this theory predicts enrichments of O, C, N, S, Ar, Kr and Xe relative to protosolar abundances due to two main factors: the settling of dust grains prior to mass loss (Mousis et al, 2018), and the possibility of photoevaporation of the gaseous envelopes of any young planet beyond 5-10 AU by nearby, newly-formed OB stars (Boss et al, 2002). While the theory of disk instability cannot yet be ruled out as the or one of the mechanism(s) of formation for Uranus and the gas giant planets in general, there are several discrepancies within models implementing this theory that are difficult to account for (Mousis et al, 2020).

The photoevaporation model states that after forming at lower temperatures farther from the Sun, ice grains would have adsorbed Ar, Kr, and Xe; subsequently migrating inward, these grains encountered increasing temperatures and those captured noble gases would have been released in this new location. These grains would essentially transport these noble gases to the regions where the Ice Giants might have formed, allowing a young Uranus to accrue supersolar noble gas enhancements, albeit smaller enrichments than those that would originate from solids containing O, C, N, and S alone (Guillot and Hueso, 2006; Mousis et al, 2020).

Lastly, the CO snowline (around 30 AU, where the temperature of the protosolar nebula was equivalent to CO’s freezing temperature) model predicts that both Uranus and Neptune were formed past the CO snowline both by solids containing C and O enrichments and gas depleted of nitrogen. Under these circumstances, Uranus’s atmosphere would contain equally depleted levels of Ar and N, but supersolar levels of Kr, He, S, and P. In comparison, C and O abundances would be very high/. Notably, this scenario would produce a D/H ratio that would be low enough to reflect observed values (e.g. Feuchtgruber et al (2013)) which, along with a higher O abundance, is often difficult for other models to produce (Ali-Dib et al, 2014).

Understanding how the planets of our Solar System formed in general is a major outstanding question, but the formation history of the Ice Giants and Uranus in particular hold the key to disentangling the exact process that resulted in the current distribution of Gas Giants, in terms of size, distance from the Sun, and bulk compositions. Due to their larger distance from the Sun than Jupiter and Saturn, Uranus and Neptune had longer available formation timescales, and so might have undergone different processes under different conditions. Additionally, understanding how Neptune-sized planets form in general is not only necessary to understand the evolution of our own Solar System, but also exoplanetary systems where Neptune-sized exoplanets are plentiful (Batalha et al, 2013). In order to constrain models attempting to balance Uranus’s total hydrogen/helium mass, their metallic composition, and the time frame of formation, the bulk abundances of several significant species must be measured accurately.

To address this question and differentiate between these formation and migration theories, it is necessary to determine the bulk abundances of a series of elements and isotopes (Mandt et al, 2020). In general, the required degree of accuracy stems from the need to compare these measurements to other known quantities throughout the Solar System. In particular, comparisons to known protosolar abundances will allow for discrimination between formation models that estimate certain bulk compositions.

Noble gases and their isotopic ratios, namely He, Ne, Ar, Kr, and Xe, are measurable only with an in situ probe. An exception is He, in that it can be partially derived from far-IR observations if temperature and the para-H₂ abundances are also constrained (Fletcher et al, 2020), albeit at a much lower level of accuracy than can be obtained in situ. Abundances of these noble gases relative to H, along with their isotopes, will aid in differentiating between the aforementioned planetary formation theories by enabling the identification of the subresevoirs of material in the protosolar nebula that fed planet formation. The He measurement should be accurate to at least $\pm$2% to enable a comparison to those made by the Galileo probe and to levels in Jupiter’s atmosphere. Similarly, Ne, Xe, Kr, and Ar should be measured at $\pm$1% to match the accuracy of those values known for the rest of the Solar System (Orton et al, 2021b).

Measurements of the bulk abundances of C, N, S, O, and P down to a pressure of at least 10 bars will inform both the nature of the objects that condensed to form the planet and the formation location relative to the carbon monoxide ice line. These elements should be obtained to accuracies of 10% or better, which is close to the current uncertainty of their protosolar abundances (Mousis et al, 2018; Atkinson et al, 2020; Hofstadter et al, 2017; Orton et al, 2021b). Some of these bulk abundances can be obtained through tropospheric measurements of CO and PH₃ (to levels of $\pm$5% for comparison to known values (Fletcher et al, 2009; Mousis et al, 2014)), which together can help constrain the deep H₂O profile, since a probe could not reach the depths below the water condensation level. Visscher and Fegley (2005) were able to use a combination of tropospheric PH₃ and CO measurements to produce upper and lower limits to the deep H₂O/H₂ abundance at Saturn through the application of thermochemical equilibrium and kinetic calculations. Similarly, Cavalié et al (2020) provided better constraints on the deep abundance of H₂O at Uranus. They reproduced CO observations with a thermochemical and diffusion model that accounted for an inhibition of convection brought on by the mean molecular mass gradient resulting from H₂O condensation at depth. To determine the nature of the reservoir of material within the protosolar nebula where Uranus was formed, the isotopic ratios of these species should also be obtained at certain accuracy levels, namely D/H and ¹⁵N/¹⁴N at $\pm$5% (Mousis et al, 2016), ³He/⁴He at $\pm$3% (at least as good as measurements made by the Galileo neutral mass spectrometer (Mahaffy et al, 1998)), and $\pm$1% for ¹⁷O/¹⁶O, ¹⁸O/¹⁶O, and ¹³C/¹²C to enable comparisons to known Solar System values for those isotopes (Mousis et al, 2018; Atkinson et al, 2020; Orton et al, 2021a, b)

The distinction between various formation models and their implications for relative abundances that would be measured by an in-situ probe was elucidated most recently by Mousis et al (2022). The clearest distinctions between different formation models lie in measurements of different enhancements or depletions of specific elements or families of elements with respect to their expected abundances in the protosolar nebula. Although schematic, Figure 2 of Mousis et al (2022) indicates enhancements and depletions resulting from the different models to be around factors of several. Therefore, the critical uncertainty controlling the differentiation between these models lies in the uncertainty of abundances in the protosolar nebula itself, which is on the order of 10% (Lodders et al, 2009). Absolute measurements more accurate than 10% would not provide additional information. For Helium in particular, a case can be made for a measurements at the $\sim$2% level that is commensurate with the uncertainty of the Galileo probe Helium Abundance Detector (HAD) (von Zahn et al, 1998). It is only the HAD measurement that decisively noted that the He abundance was below what one expected in a protosolar nebula, and thus subject to dissolution in the deeper atmosphere. Current interior models (e.g., Guillot (2005); Lambrechts et al (2014); Helled et al (2020)) do not expect this to be the case for Uranus or Neptune, and so it is important to determine its abundance to the same level of uncertainty as for the Galileo mission.

Uranus’s heat budget, or the balance between solar insolation and radiative heat loss, is enigmatic among Gas Giants as it appears to be half that of Neptune’s despite being closer to the Sun and receiving more radiant solar energy (Pearl et al, 1990; Pearl and Conrath, 1991). Uranus’s atmosphere is the medium through which residual gravitational potential energy from planetary formation is released, and thus, can heavily influence the planet’s heat budget and its thermal evolution. In general, the processes that transport this internal heat energy through gas giant atmospheres are cumulus convection and radiation (Li and Ingersoll, 2015); however, some other process might also be occurring in Uranus’s atmosphere that either inhibits this type of heat flow or causes atypical patterns of heat transport altogether.

Several hypotheses have been put forward to explain Uranus’s apparent low internal heat flux and luminosity. Some deep compositional gradient and resultant thermal boundary layer, possibly initiated by a giant impact early in Uranus’s history (Kegerreis et al, 2018), might be responsible for trapping heat deep in the atmosphere and inhibiting its vertical transport (Nettelmann et al, 2016; Vazan and Helled, 2020; Scheibe et al, 2021). For example, the presence of even a thin conductive layer in Uranus’s interior has been predicted to dramatically affect planetary cooling (Scheibe et al, 2021). Another possibility is simply the fact that various potential issues with the Voyager measurements of Uranus’s luminosity misled estimates of the amount of internal heat. Possible issues include limitations on phase angle coverage, possible calibration issues for Voyager instruments (Li et al, 2018), and/or temporal changes in reflectance (Lockwood, 2019) since those measurements were made. Or, we might simply be observing Uranus during a transient, quiet period of heat release relative to Neptune; a discussion of the meteorological causes of such a quiescent period can be found in Section 2.3.1. Li et al (2015) found that Saturn’s globally-averaged radiant power to space increased by $\sim$2% during its 2010–2011 Great White Spot outbreak, which renders it possible that a quiescent and less-stormy Uranus might have appeared to have a low outgoing heat flux during the Voyager epoch.

To test the above hypotheses, the temperature/pressure profile should be measured to identify any regions with sub- or super-adiabatic lapse rates. Identifying deviations from adiabatic lapse rates can reveal how and to what degree internal heat is being transported farther down in the atmosphere. To accurately measure these temperature/pressure profiles, pressure measured with an accuracy of $\pm$1% and temperature to $\pm$1 K will be required (Orton et al, 2021a). These precision requirements arise from the need to differentiate between different lapse rates and degrees of atmospheric stability as observed by Voyager 2 and predicted by works such as Guillot (1995), Leconte et al (2017), and Cavalié et al (2020) in the middle and deep troposphere and Orton et al (2014) and Roman et al (2020) in the upper troposphere and stratosphere, as well as the need to locate any radiaitve boundary layers. Using radio occultation measurements, Voyager 2 measured temperature down to $\sim$2 bars with an accuracy of 1 K and at best 1% accuracy for pressure (Lindal et al, 1987), and so a comparable accuracy is necessary to understand how Uranus’s temperature profile has changed since 1986. A sampling frequency of 1 Hz will allow for multiple ($\sim$10, depending on the probe’s descent rate) measurements per scale height (e.g., Sayanagi et al (2020)).

Additionally, the vertical profile of condensable species abundances, especially those containing C, N, S, O, and P, particularly CH₄ and H₂S (H₂O is unlikely to be feasible), ideally at the levels described in Section LABEL:sec:2.1.3, but for these purposes 10-20% will suffice since matching well-constrained Solar System or protosolar values is not required. Measuring condensable abundances multiple times per scale height is the minimum sampling frequency needed to address this science question and contextualize the measured lapse rate, but higher sampling frequencies are highly desirable to provide better constraints to extant model predictions (Ferri et al, 2020). Such profiles will also be valuable for identifying any potential sources or sinks of those species that might be populating the compositional/thermal boundary layer (e.g., Cavalié et al (2020)). Additionally, a microwave radiometer onboard the orbiter would be better able to access pressures and reservoirs of gases that the probe cannot.

Also necessary is the vertical profile of horizontal winds to $\pm$10 m/s, which can be obtained through Doppler measurements of the probe’s position as it descends through the atmosphere (Sayanagi et al, 2020). These measurements will help reveal the way in which wind shears with depth, what sort of waves it might encounter, and will allow for linking wind shear and the horizontal temperature gradient through the thermal wind relationship (Orton et al, 2021b).

The degree of vertical convection, and therefore the convective capability at cloud condensation pressures, is critical to understanding the way in which atmospheric stability affects heat transfer (Banfield et al, 2005). This information can be obtained through measurements of the speed of sound to accuracies of $\pm$1% in a given region of the atmosphere, which can then can be used to derive the ortho-para H₂ fraction at that location (Orton et al, 2021a).

To identify the direction of heat flow in the atmosphere and any resulting differences in buoyancy, the net atmospheric radiative energy balance should be measured by finding the difference between downwelling solar energy and upwelling thermal infrared energy (Atkinson et al, 2020). In order to constrain those two contributors to the radiative energy flux, the altitude profile of thermal infrared light and the amount of absorbed visible sunlight should be determined. The amount of absorbed sunlight can be found once the planet’s Bond albedo is characterized, which requires remote sensing measurements from the orbiter. Visible flux, in both directions, should be measured from 0.4–5 $\mu$m to fully cover the contribution from solar flux, and the IR flux should be measured over 4–50 $\mu$m, which is much wider than similar requirements for Neptune (Hofstadter et al, 2017) to account for Uranus’s possibly much lower themral emissoin. Both fluxes should be acquired with a spectral resolution ($\lambda/\Delta\lambda$) of 0.1–100 over their respective wavelength ranges, a flux resolution of $\sim$0.5 W/m^-2, and an accuracy of $\pm$1% (Hofstadter et al, 2017; Orton et al, 2021a).

The dominant mechanism of vertical heat transport in the troposphere, aside from incoming solar radiation and outgoing thermal heat from the interior, is cumulus convection driven by the condensation of CH₄, H₂O, NH₄SH, and possibly NH₃ and/or H₂S (Hueso et al, 2020). This process likely plays an important role in the energy budget of the atmosphere and in storm generation (Sromovsky and Fry, 2005). Could this convection, and its associated release of latent heat and convective available potential energy (CAPE), be behind the episodic storms observed in Uranus’s atmosphere?

In giant planets’ hydrogen-dominated atmospheres, condensable species have a significantly higher molecular weight than the background H₂-He atmosphere. Consequently, the presence of heavier vapor made up of those condensables can suppress cumulus convection for prolonged periods and may allow for the build up of an enormous amount of convective available potential energy (CAPE) before convection is triggered (Li and Ingersoll, 2015). This process is thought to be behind the roughly 30-year cycle on Saturn’s Great White Spot outbreaks on Saturn (Sánchez-Lavega et al, 2018; Li and Ingersoll, 2015). When sufficient CAPE is accumulated and latent heat release becomes capable of overcoming the stabilizing heavy vapor, an episodic cumulus storm may erupt, enhancing the heat flux to space. On Uranus, the interval between such episodic storms may be decades or centuries. See Section 2.6.1 for a discussion of periodic and potentially seasonal cloud features that have been observed on Uranus.

During a quiescent period, storm activity may be very low and the resulting heat transport to space may fall to levels comparable to the incoming solar flux. Alternatively, or perhaps in conjunction with convective suppression from stabilizing layers of high vapor abundances, more deeply seated and less efficient radiative heat-transporting layers may exist on Uranus. The effect of such a quiescent period may be hiding a warmer-than-expected interior, while simultaneously appearing as a sluggish and bland upper troposphere above a much deeper and more convectively active layer. For more details on moist convection in gas giant atmospheres and associated phenomena, see Palotai et al (2022).

Uranus and Neptune appear to harbor large reservoirs of condensable species (e.g., CH₄, H₂O). When these vapors condense in an upwelling air parcel, they release latent heat, with H₂O being particularly energetic. However, water clouds are situated deep in the troposphere, perhaps at 50–200 bar, depending on the global abundance of O (Cavalié et al, 2020). CH₄ clouds condense much higher, around 1.3 bar (Irwin et al, 2022), and can be easily sampled by an entry probe assuming the probe passes through these ephemeral and localized clouds. The purported globally distributed H₂S cloud deck around 3–6-bar will also be accessible to the entry probe; most probe designs target at least the 10 bar pressure level (Mousis et al, 2014, 2018; Orton et al, 2021a; Simon et al, 2020). Thermochemical equilibrium models suggest that a deeper NH₄SH cloud layer exists around 30–40 bar (Weidenschilling and Lewis, 1973; Atreya and Wong, 2005), but this may be too deep for the probe to measure within the constraints of battery power, telecommunication restrictions, or increasing radio wavelength opacity at and below that pressure level due to the presence of NH₃. However, sampling the upper CH₄ and H₂S cloud decks may be sufficient to constrain some of the circulation hypotheses, aerosol properties, and will provide invaluable data on the zonal wind profile.

Moist convection may also play a role in forming the dark anticyclonic spots, more commonly found on Neptune (e.g., Smith et al (1989)), but these have also been more rarely observed on Uranus (Hammel et al, 2009). Alternatively, these dark spots may form as a result of baroclinic or barotropic instabilities. These vortices are usually associated with bright companion clouds that sometimes exhibit rapidly changing morphologies, perhaps linked to convective activity. These spots and their companion clouds are much more rare on Uranus, which suggests they are linked to a more vigorously convective period, such as may currently exist on Neptune. It is not yet known if their dark appearance (in blue wavelengths; nearly undetectable in red wavelengths) are a result of lower aerosol abundance (cloud-clearing) in the H₂S cloud deck, or a change in the physical character of the aerosols there (Hadland et al, 2020). However a recent study using the Multi Unit Spectroscopic Explorer instrument (MUSE) of Neptune’s NDS-2018 dark spot (Irwin et al, 2023) ruled out the cloud-clearing hypothesis. It remains to be seen if Uranus’s dark spots are also a result of different aerosol properties like Neptune’s but, given their similar atmospheric compositions, it is a likely scenario. While it is extremely unlikely a probe will descend into a dark spot, measuring representative aerosol populations at the altitudes these dark spots are observed will help constrain radiative transfer models.

Establishing a vertical profile of the temperature and composition of the atmosphere as a function of probe altitude is mandatory to determine the stability of the atmosphere against convection. Temperature and pressure should be captured at accuracies of $\pm$1 K and $\pm$1%, as previously stated. Condensables, such as CH₄, H₂S, and NH₃, should be measured to accuracies at approximately $\pm$20%, down to at least 5 bars, and at least multiple times per scale height (Sayanagi et al, 2020). These measurements will establish what parts of the troposphere, if any, contain super-adiabatic lapse rates, which can only exist due to the convectively-suppressing molecular weight gradient effect from condensable vapors of a sufficient concentration (i.e., specific humidity) (Lian and Showman, 2010). Additionally, measurements of the ortho-para H₂ abundance as a function of altitude can be used to quantify the degree of vertical mixing (Banfield et al, 2005). Determining vertical profiles of these parameters will also help determine the amount of convective potential in the atmosphere and quantify the atmospheric stability.

In a paper detailing the science case and design for a small probe to be carried along with a larger and more heavily equipped probe, Sayanagi et al (2020) recommends an alternative variable accuracy for the temperature measurement: $\pm$5 K between 1 and 10 mb (stratosphere), $\pm$2 K between 10 and 100 mb (stratosphere to tropopause), and $\pm$0.1 K at greater pressures to meaningfully distinguish lapse rates. For an atmospheric descent that takes 2,235 seconds between 100 mbar and 10 bar, a sampling rate of 1 Hz for both temperature and pressure measurements will ensure that at least 10 measurements are performed every scale height.

Doppler wind measurements, through the use of ultrastable oscillators on the probe and carrier spacecraft, can in principle detect the presence of atmospheric waves (Seiff et al, 1997). A Doppler wind experiment will establish a vertical profile of the horizontal winds (e.g. Atkinson et al (1996, 1997, 1998)). These winds and waves, measured to an accuracy of $\pm$10 m/s (Orton et al, 2021a), provide clues to atmospheric dynamical processes occurring at the probe entry site and are important in determining a more accurate measurement of zonal jet speeds. Acceleration measurements suggested by Sayanagi et al (2020) need to span a large range, from 0.01 m/s² for a reconstruction of horizontal winds of $\pm$10 m/s, and up to 300 G during entry. A sampling frequency at greater than 50 Hz is specified in their report.

Above the stratosphere, Uranus’s thermosphere was measured by Voyager 2 to be warmer than expected (Broadfoot et al, 1986), but has been cooling at a rate of 8 K/yr since 1997 (Melin et al, 2019, 2020). Possible explanations for the thermosphere’s thermal evolution include magnetic effects, such as auroral heating and the effects of daily magnetic field reconnections (e.g. Cao and Paty (2017)), or some sort of atmospheric process.

Despite receiving lower levels of insolation than Jupiter and Saturn, Uranus has relatively active photochemical processes taking place in the upper atmosphere. In the stratosphere and above, lofted and UV-dissociated CH₄ molecules generate an abundance of complex hydrocarbons and photochemical products (Moses and Poppe, 2017; Moses et al, 2018, 2020). These products can effectively absorb solar energy from above and inhibit the release of heat from deeper in the atmosphere. Uranus’s CH₄ homopause is lower than Neptune’s, at regions where diffusion timescales are large (Moses et al, 2018). Meridional stratospheric circulation is also an unlikely candidate for the cause of this uplifting, since it also appears to be too lethargic (Flasar et al, 1987). It is therefore likely that some more deeply-rooted mechanism might be lofting that CH₄ to levels where it can be affected by solar radiation. It is also possible that seasonal changes to the eddy diffusion coefficient, driven by rising temperatures, increased the amount of hydrocarbons in the upper atmosphere since Voyager visited (Herbert et al, 1987; Bishop et al, 1990) and in turn increased that region’s cooling efficiency, thereby allowing the thermosphere to cool (Melin et al, 2020).

Identifying the dynamical mechanisms that influence the cooling and radiative processes in the upper atmosphere is key to narrowing down the possible causes of the enigmatic thermal evolution of Uranus’s thermosphere. First, measuring and characterizing the sources, sinks, and behavior of the products of CH₄ photolysis is necessary in order to use them as dynamical tracers. Additionally, CO has been found in the stratosphere above the 100-mbar level (Cavalié et al, 2014) and could be used as a dynamical tracer, as well as HCN if it is detected in Uranus’s upper atmosphere, as it has been in Jupiter’s (Moreno et al, 2003; Cavalié et al, 2023).

A probe can help disentangle the magnetic and atmospheric influences on the thermal evolution of the upper atmosphere by complementing remote sensing measurements of the top of the atmosphere. In particular, a probe can explore deeper dynamical effects and measure species abundances below $\sim$100 mbar (the point at which the probe will be stable enough to make accurate measurements; e.g., Sayanagi et al (2020)). Wind speeds and cloud motion, measured within $\pm$5 m/s to compare to dynamical models (Orton et al, 2021a), will both be important characteristics to measure as the probe descends through the atmosphere, and useful for assessing the degree of circulation and uplift in regions below the stratosphere. Temperature/pressure profiles, with the same constraints as previously discussed, will be important for directly assessing the nature of heat transport in the region where the probe is dropped. The vertical composition profile of both elemental abundances and condensable species, especially CH₄, will be necessary to within $\pm$10% (Orton et al, 2021a); again, with as many measurements of these values per scale height as data volume allows.

If possible at altitudes at and above the 100-mbar level, identifying the particle optical properties, size distributions, and number and mass densities, as well as the opacity, shapes, and composition of any photochemical hazes in the stratosphere will be critical to constrain the sources and sinks of the efficiently-radiating hydrocarbons (Moses et al, 2020). However, the probe will likely be unable to accomplish these measurements in the stratosphere due to the timing of parachute deployment. Therefore, the orbiter should be equipped to take similar or complementary measurements to characterize the photochemical products of CH₄ photolysis.

The probe’s suite of instruments will primarily sound Uranus’s tropsophere. This region of the atmosphere harbors zonal wind patterns observed for both Ice Giants that are markedly different in three main ways from those observed on the Gas Giants: 1) There is a single broad equatorial jet with retrograde motion relative to the planet’s rotation; 2) There is only one prograde mid-latitude jet per hemisphere; 3) The width of these Ice Giant jets span up to several tens of degrees in latitude with velocities up to $\sim$700 m/s, making them both wider and faster than jets found on Jupiter and Saturn (Sromovsky and Fry, 2005; Ingersoll, 1990). Uranus in particular appears to contain belts and zones of darker and brighter reflectivity (e.g. Sromovsky et al (2012)), but these are not closely correlated with the position or latitudinal gradient of the jets, as is true for Jupiter (e.g., García-Melendo and Sánchez-Lavega (2001)). Lastly, the jets on the Ice Giants do not penetrate as deep into the interior as they do on the Gas Giants (1,000 km at the Ice Giants vs. 3,000 km at Jupiter and 9,000 km at Saturn (Kaspi et al, 2013, 2020)). The differences between the jets of the Ice Giants and those of the Gas Giants, particularly on Uranus under the influence of its extreme seasonal forcing, the cause of those differences, and how they influence other processes of atmospheric circulation are open questions.

With fewer cloud features available for tracking, combined with fewer observations, and large distances from Earth, we do not have a rich spatial and temporal data set to compare to the closer and more frequently visited Gas Giants. From a dynamics perspective, additional questions arise about Uranus’s jets: 1) Do they violate the Rayleigh-Kuo and Charney-Stern stability criteria by large margins as they appear to do on Jupiter? 2) What is the magnitude of the eddy flux momentum—which powers the jets on the Gas Giants—into or out of the jets (Salyk et al, 2006; Del Genio et al, 2007)? Is this flux driven by moist convection as is thought on Jupiter and Saturn (Ingersoll et al, 2000)? We must keep in mind that zonal wind measurements on Uranus do not even span a complete Uranian year, therefore 3) Does this flux vary greatly over time, and if so, how may that affect the zonal and meridional wind circulations?

The lower number of jets on the Ice Giants compared to the Gas Giants ($\sim$16 vs. 10 hours) are likely produced by the combination of the smaller planetary radii and slower rotation rates. In short, the Rhines length, which is a quantity that provides an estimate of zonal wind width, is much smaller than the Gas Giants’ planetary radii (resulting in more jets), but comparable to the Ice Giants’ planetary radii (Cho and Polvani, 1996). However, this does not explain why the equatorial jet is retrograde-directed, but may be a consequence of reduced convective activity near the equator, unlike that seen on the Gas Giants (Lian and Showman, 2010).

The zonal jets are expected to decay upwards into the upper troposphere and, at much deeper levels below the troposphere, downwards toward the interior, but neither the upper nor lower decay profile is well constrained (Conrath et al, 1990; Kaspi et al, 2013; Fletcher, 2022). However, a recent ground-based observation of Neptune from the Atacama Large Millimeter/submillimeter Array (ALMA) (Carrión-González et al, 2023), found retrograde stratospheric jets in the range of 170 m/s to 190 m/s near the equator at pressures of 2 mbar and 0.4 mbar, respectively. These wind speeds are in agreement with stellar occultation measurements and expectations from the thermal wind equation. Given the similarity between Uranus’s and Neptune’s tropospheric zonal jets, we would expect to find similar stratospheric zonal jets at Uranus. Measuring vertical wind shear in the upper atmosphere will provide much needed altitude constraints for tracking features from remote sensing observations. Wind shear will also provide an indication of the horizontal temperature gradient via the thermal wind relationship (Holton, 2004). A probe’s Doppler wind experiment will be able to provide good constraints on the wind profile in the stratosphere and upper troposphere, although the probe will long cease functioning before reaching either the water cloud layer or the depth where zonal jets are expected to decay (Kaspi et al, 2013).

Interestingly, Uranus’s equatorial jet around equinox in 2012 revealed a braid-like wave feature that appears unique in the Solar System (Sromovsky et al, 2012). Slightly north of Jupiter’s equator, cloud features are often observed that may be manifestations of an equatorial Rossby wave. We do not know what type of wave might explain Uranus’s equatorial braid, but it did not survive for more than a few years. This feature, and our experience with the Galileo probe entering a hot spot on Jupiter devoid of water signatures, needs to be considered when studying Uranus’s equatorial dynamics and data obtained from a probe entering this region. Furthermore, several wave-like phenomenon have been observed on Jupiter which provide information about dynamical stability of the region and the processes that drive atmospheric flows (Simon et al, 2018). Similarly, the Juno spacecraft observed a plethora of small scale wave-like phenomenon on Jupiter (Orton et al, 2020), providing details about local dynamics. Understanding the nature and distribution of such features on Uranus will enable key insights into both the local and global dynamical structure of the Uranian atmosphere, and formation of dynamical instabilities.

From efforts to reconcile observed temperatures, opacities, zonal wind profiles, and the results of numerical modeling, a new hypothesis has emerged to explain these varied characteristics in Jupiter’s atmosphere: the presence of a double-stacked set of overturning cells in the troposphere (Ingersoll et al, 2000; Showman and de Pater, 2005; Fletcher et al, 2020; Yamazaki et al, 2005; Zuchowski et al, 2009). In the upper troposphere ($\sim$0.1–1 bar), large-scale meridional circulation suggests upwelling occurs in the midlatitudes with downwelling occurring over the equator and poles. Deeper, a mid-troposphere cell (1 bar to an undetermined depth) exists and is characterized by counter-rotation to that of the upper tropospheric cell. Such stacked, counter-rotating cells have been invoked to explain seemingly conflicting results of observations in Jupiter’s belts and zones. The Juno spacecraft has offered evidence to support the double-stacked hypothesis on Jupiter (Fletcher et al, 2021). However, a dynamically consistent picture that connects aerosol and temperature observations together with the zonal-mean circulation is currently incomplete for the Ice Giants. An orbiter will be necessary to determine if the double-stacked cell hypothesis is valid at Uranus, with any degeneracy in solutions provided by remote sensing potentially being resolved with probe measurements.

The measurement requirements to address the circulatory and zonal dynamics of Uranus’s troposphere are identical to those in Section 2.3.2. The one additional requirement is a measurement of ortho-para H₂ fraction via speed of sound measurements at accuracies of $\pm$1%. In this case, the ortho-para H₂ fraction will act as a dynamical tracer in the upper troposphere. The ortho-para H₂ ratio is higher at higher temperatures deeper in the atmosphere, and once a parcel of air from this region containing that ratio convects to the upper troposphere, the conversion to a lower ortho-para H₂ fraction takes years due to colder temperatures higher up. Therefore, we can use the ortho-para H₂ ratio at altitude as a measure of vertical motion (Banfield et al, 2005).

Neptune’s and Uranus’s circulation seem to be broadly similar despite Uranus’s 98°obliquity and much higher degree of insolation at its poles relative to its equator. Uranus’s deep atmosphere has displayed behavior consistent with seasonal changes after the southern solstice in 1985 (Hofstadter and Butler, 2003), and more recent Very Large Array observations have shown potential seasonal changes in the brightness temperature of the polar cyclone at Uranus’s north pole (Akins et al, 2023). Bright and short-lived cloud features have been occasionally observed at the top of the atmosphere (e.g. Hammel et al (2005)), and while it is yet unclear what causes these events and if those effects are seasonal in nature, they might be linked to cumulus convective outbreaks or local condensation resulting from vortices causing vertical displacement in the atmosphere (Sromovsky et al, 2012). In general, cloud and apparent convective activity increased as Uranus approached and passed the 2007 equinox, until 2014. After 2014, observable convective activity appeared to cease, with the exception of occasional white clouds near the polar hood (Hueso et al, 2020).

Remote-sensing instruments onboard an orbiter are likely best-suited for addressing this question, in that they can obtain observations over much more expansive temporal, spatial, and longitudinal ranges than a probe. However, a properly-equipped in situ probe would be able to greatly enhance our ability to answer this question by directly probing the atmospheric structure and processes deeper down. Particularly, measuring the lapse rate and how it changes with altitude offers insight into the degree of vertical mixing and convective strength as a function of depth. How the upper and lower regions of Uranus’s atmosphere influence one another is essential to understanding the planet’s energy budget and the interplay between seasonal and non-seasonal forcings.

Specifically, several questions remain on the seasonal dynamics of the Uranian atmosphere. (1) How does the global circulation change on the seasonal timescale? Ground-based observations have shown that cloud and storm formation is tied to specific seasons (de Pater et al, 2015), but is this a result of global upwelling and downwelling or localized orographic effects from polar vortices? (2) How does the thermal flux change on the seasonal timescale? We require more measurements of the thermal flux on Uranus over multiple seasons in order to understand this variability. (3) How do aerosol properties change seasonally? Do the haze particles vary significantly with seasonal forcing? It is vital that we answer these questions so as to generate a model that unifies the various processes that drive the atmosphere (i.e., dynamics, microphysics, thermal and mechanical forcing, etc.).

A probe capable of measuring changes in local and planet-wide circulation, especially when compared to Voyager 2 measurements, would dramatically increase the science return of a mission seeking to address this question. As the probe descends through the atmosphere, measurements of the vertical and horizontal wind profiles will provide vastly improved constraints for 3D circulation models of Uranus’s atmosphere.

Global circulation modeling studies would greatly benefit from even a single in situ data point, specifically in constraining the vertical profile of temperature, wind speeds and aerosol density, all to accuracy levels previously described. These properties will prove critical to constraining current dynamical models (e.g., Hammel et al, 2009; Helled et al, 2010; Kaspi et al, 2013; LeBeau et al, 2020), particularly with relation to the processes that drive the deeper atmospheres.

Additionally, in order to holistically determine atmospheric composition, we require measurements of the scattering cross section and spectral identification of haze aerosol species in the upper atmosphere as much as the probe is able below the lowest pressure limit of 100 mbar, particularly as they pertain to seasonal concentrations of aerosols and hydrocarbons. The ratio of trace species to the bulk abundances (e.g., of C, N, O and S) enable estimates on the vertical mixing length and timescale, as well as the degree of vertical circulation.

Furthermore, arriving at Uranus before or during the 2049 equinox will provide the best opportunity to measure the atmosphere when the planet is at the largest disparity between measurement epochs by dedicated spacecraft. Voyager 2 arrived at Uranus during solstice and took relatively limited measurements due to the nature of its quick flyby; orbiting and probing Uranus during equinox will enable a direct comparison of the two spacecrafts’ datasets at opposite seasonal epochs. Arriving during a time frame leading up to and/or after an equinox enables a comparative planetology experiment between Uranus and other bodies that are particular active at equinox, such as Titan (Coustenis et al, 2020).