Astronomy

How did UV from the earliest stars 'alter the state of the 21 cm line' such that it shows up in CMB today?

How did UV from the earliest stars 'alter the state of the 21 cm line' such that it shows up in CMB today?


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In this question I discuss the recent (open access) paper in Nature An absorption profile centred at 78 megahertz in the sky-averaged spectrum at length. The abstract begins:

After stars formed in the early Universe, their ultraviolet light is expected, eventually, to have penetrated the primordial hydrogen gas and altered the excitation state of its 21-centimetre hyperfine line. This alteration would cause the gas to absorb photons from the cosmic microwave background, producing a spectral distortion that should be observable today at radio frequencies of less than 200 megahertz(1).

where reference 1 is Jonathan Pritchard and Abraham Loeb (2012) 21 cm cosmology in the 21st century. From there I've found Leonid Chuzhoy and Zheng Zheng (2007) Radiative Transfer Effect on Ultraviolet Pumping of the 21 cm Line in the High-Redshift Universe.

I understand some basics about the hyperfine transition in hydrogen, and that the "spin temperature" of a gas in space can differ from the temperature of other partitions if it's being pumped, but these papers are more than a bit hard to read.

Is there a simple way to explain the basics behind how the exposure of hydrogen to the UV light produced in early stars would cause the "blip" in the (now red-shifted) 21 cm part of the Cosmic Microwave Background radiation spectrum we see today?


Guiseppe Rossi's answer is excellent; I just want to add why the mentioned UV radiation modifies the spectrum of the background radiation.

(one word took the other, and it became a rather long comment.)

The hyperfine level

Neutral hydrogen in its ground state can be in two different configurations; either the proton and the electron may have parallel spins ($uparrowuparrow$), or they may have antiparallel spins ($uparrowdownarrow$). When the spins are parallel, the atom has a slightly higher energy than when they're antiparallel. The atoms "wants" to make a spin flip to the lower energy configuration$^dagger$, and will eventually do so, but since the line is forbidden, the lifetime of the parallel state is of the order $10^7,mathrm{yr}$.

The relative population of the states is given by the Boltzmann distribution $$ egin{array}{rcl} frac{n_1}{n_0} & = & frac{g_1}{g_0} , e^{-Delta E , / , k_mathrm{B} T_S} & = & 3 , e^{-0.068,mathrm{K} , / , T_S}label{a} ag{1}, end{array} $$ where subscripts 1 and 0 denote the $uparrowuparrow$ and $uparrowdownarrow$ states, respectively, $n$ is the density, $g$ is the statistical weights (with $g_0,g_1 = 1,3$), $Delta E = 5.9 imes10^{-6},mathrm{eV}$ is the energy difference of the states, $k_mathrm{B}$ is the Boltzmann constant, and $T_S$ is the spin temperature, which I think is better thought of as "a number that describes the relative populations" than an actual temperature.

Departure from equilibrium

In thermal equilibrium, the spin temperature is equal to the "real", kinetic temperature. Just after decoupling of the radiation from matter at a redshift of $zsimeq1100$, the gas and the photons share the same energy, and since $Tgg 1$, we have that $n_1/n_0 simeq 3$. But when the first stars begin to shine, they produce massive amounts of hard UV radiation which ionizes their surrounding medium. The ionized gas quickly recombines (in the beginning, at least), with $sim2/3$ of the recombinations resulting in the emission of a Lyman $alpha$ photon, i.e. a photon with an energy corresponding to the energy difference between the first excited state (one of the three $2P$ states) and the ground state (the $1S$ state) of the hydrogen atom (10.2 eV).

The Ly$alpha$ photons scatter multiple times on the neutral hydrogen. Each scattering excites an atom from $1S ightarrow 2P$, which subsequently de-excites and emits an Ly$alpha$ photon in another direction. But since the energy difference between the $2P$ and the $1S$ state is a million times larger than between the hyperfine states, there is equal chance of ending in the $uparrowuparrow$ and the $uparrowdownarrow$ state.

That is, $n_1/n_0$ is no longer $simeq 3$, but is driven toward $sim 1$. This is the Wouthuysen-Field effect that Guiseppe Rossi mentions; from eq. ef{a}, you see that this corresponds to a much smaller spin temperature, and thus the factor that Guiseppe Rossi mentions becomes negative. The full equation describing the brightness (or, equivalently, the flux received) as a function of redshift can be written (e.g. Zaldarriaga et al. 2004) $$ T(z) = 23,mathrm{mK} , frac{T_S - T_mathrm{CMB}}{T_S} , (1+delta) x_mathrm{HI}(z) frac{Omega_mathrm{b}h^2}{0.02} left( frac{0.15}{Omega_mathrm{m}h^2} , frac{1+z}{10} ight)^{1/2}label{b} ag{2} $$ and when the $(T_S - T_mathrm{CMB}),/,T_S$ factor is negative, you will get an absorption line.

(In eq. ef{b}, $delta$, $x_mathrm{HI}(z)$, $Omega_mathrm{b}$, $Omega_mathrm{m}$, and $h$, are the local overdensity, the neutral fraction of hydrogen, the baryon and matter density parameter, and the dimensionless (reduced) Hubble constant, respectively, but this is of less importance.)

Since the observed absorption line (Bowman et al. 2018) starts to drop around an observed frequency of $ u_mathrm{obs} = 65 ext{-}70,mathrm{MHz}$, and since the rest frequency of the hyperfine line is $ u_mathrm{rest} = 1420,mathrm{MHz}$, this means that the first stars appeared around a redshift of $z = u_mathrm{rest}/ u_mathrm{obs} - 1 simeq 20$, corresponding to an age of the Universe of $sim 180,mathrm{Myr}$ (i.e. million years - the largest absorption is reached at $zsimeq 17$, or $tsimeq 200,mathrm{Myr}$).

Now the big question is, according to eq. ef{b} the dip should be of the order a few tens of mK, but is in fact roughly 0.5 K, i.e. an order of magnitude larger. One possible mechanism that could produce this effect is coupling of the gas with dark matter, something which is not usually considered possible but could happen if the dark matter particle has a very small charge (Barkana et al. 2018).

Time evolution of the 21 cm signal

The figure below (from a great review by Pritchard & Loeb 2012) shows how the 21 cm signal evolves with time. The dip discussed in this answer is the orange and red part.

$^dagger$An analogy would be two magnets aligned parallel to each other with north in the same direction, preferring to flip around, but note Ken G's comment below; the transition doesn't necessarily involve a spin flip, and the analogy is not to be taken literally, since parallel magnets are alike, whereas parallel electrons/protons have opposite charges.


The Intergalactic medium at the relevant redshift is made of neutral hydrogen. What we can measure is the brightness temperature (the temperature that the IGM would have if it emitted as a blackbody) relative to the CMB. This quantity depends crucially on the following expression:

$frac{T_S - T_{CMB}}{T_S}$

where $T_S$ is called spin temperature.

The spin temperature is just a measure of the ratio between the number of hydrogen atoms in the first excited hyperfine state (spin parallel) and ground state (spin antiparallel). It turns out that the spin temperature can be modified in three ways: one of them is Lyman resonant scattering (the other two are collisional coupling and scattering of CMB photons). Using atomic physics it's not difficult to see that $T_S$ is a weighted mean of the temperature of the CMB and gas temperature.

UV photons can change the spin temperature because a hydrogen atom in the lowest state, n=1 (if you are familiar with chemistry, a s state) with antiparallel spins, can absorb an UV photon and jump in a p (n=2) state, and then fall again in a s state but with parallel spins. This mechanism is known as Wouthuysen-Field coupling.

Before the emission takes place, the gas is at the same temperature of the CMB and both are equal to spin temperature, so the relative brightness temperature is zero. When Lyman alpha emission begins the spin temperature decreases resulting in the observed absorption peak. At some point the peak stops because the first stars emit X-rays and the gas becomes hotter than the background radiation.


The anomalous 21-cm absorption at high redshifts

The EDGES collaboration has reported the detection of a global 21-cm signal with a plateau centered at 76 MHz (i.e., redshift 17.2), with an amplitude of (500^<+200>_<-500>) mK. This anomalous measurement does not comport with standard cosmology, which can only accommodate an amplitude (lesssim 230) mK. Nevertheless, the line profile’s redshift range ( (15lesssim zlesssim 20) ) suggests a possible link to Pop III star formation and an implied evolution out of the ‘dark ages.’ Given this tension with the standard model, we here examine whether the observed 21-cm signal is instead consistent with the results of recent modeling based on the alternative Friedmann–Lemaître–Robertson–Walker cosmology known as the (R_>=ct) universe, showing that – in this model – the CMB radiation might have been rethermalized by dust ejected into the IGM by the first-generation stars at redshift (zsim 16) . We find that the requirements for this process to have occurred would have self-consistently established an equilibrium spin temperature (T_>approx 3.4) K in the neutral hydrogen, via the irradiation of the IGM by deep penetrating X-rays emitted at the termination shocks of Pop III supernova remnants. Such a dust scenario has been strongly ruled out for the standard model, so the spin temperature ( (sim 3.3) K) inferred from the 21-cm absorption feature appears to be much more consistent with the (R_>=ct) profile than that implied by (Lambda ) CDM, for which adiabatic cooling would have established a spin temperature (T_mathrm(z=17.2)sim 6) K.


Abstract

[1] Solar X-ray and UV radiation (0.1–320 nm) received at Earth's surface is an important aspect of the circumstances under which life formed on Earth. The quantity that is received depends on two main variables: the emission of radiation by the young Sun and its extinction through absorption and scattering by the Earth's early atmosphere. The spectrum emitted by the Sun when life formed, between 4 and 3.5 Ga, was modeled here, including the effects of flares and activity cycles, using a solar-like star that has the same age now as the Sun had 4–3.5 Ga. Atmospheric extinction was calculated using the Beer-Lambert law, assuming several density profiles for the atmosphere of the Archean Earth. We found that almost all radiation with a wavelength shorter than 200 nm is attenuated effectively, even by very tenuous atmospheres. Longer-wavelength radiation is progressively less well attenuated, and its extinction is more sensitive to atmospheric composition. Minor atmospheric components, such as methane, ozone, water vapor, etc., have only negligible effects, but changes in CO2 concentration can cause large differences in surface flux. Differences due to variability in solar emission are small compared to this. In all cases surface radiation levels on the Archean Earth were several orders of magnitude higher in the 200–300 nm wavelength range than current levels in this range. That means that any form of life that might have been present at Earth's surface 4–3.5 Ga must have been exposed to much higher quantities of damaging radiation than at present.


Evidence for Diffuse Gas at High z

CMB Fluctuations as a Probe of the Ionization Epoch.

If the intergalactic medium were suddenly reionized at a redshift z, then the optical depth to Thomson scattering back to zi(≫1) would be 2 (the generalization to a more realistic scenario of gradual reionization is straightforward). Even when this optical depth is far below unity, the ionized gas constitutes a “fog” that attenuates the fluctuations imprinted at the recombination era the photons that are scattered at <zi then manifest a different pattern of fluctuations, characteristically on larger angular scales. This optical depth is consequently one of the parameters that can, in principle, be determined from CMB anisotropy measurements. It is feasible to detect a value as small as 0.1—polarization measurements may allow even greater precision, because the scattered component would imprint polarization on angular scales of a few degrees, which would be absent from the Sachs–Wolfe fluctuations on that angular scale originating at trec.

Twenty-one-centimeter Emission, Absorption, and Tomography.

The 21-cm line of HI at redshift z would contribute to the background spectrum at a wavelength of 21(1 + z) cm. This contribution amounts to a brightness temperature of order 0.05 (1 + z)½. This is very small compared with the 2.7 K of the CMB and smaller still compared with the nonthermal background, which swamps the CMB, even at high galactic latitudes, at the long wavelengths where high-z HI should show up. Nonetheless, inhomogeneities in the HI may be detectable, because they would give rise not only to angular fluctuations but also to spectral structure (5, 6). If the same strip of sky were scanned at two radio frequencies differing by (for example) 1 MHz, the temperature fluctuations due to the CMB itself, to galactic thermal and synchrotron backgrounds, and to discrete sources would track each other closely. Contrariwise, there would be no correlation between the 21-cm contributions, because the two frequencies would be probing “shells” in redshift space whose radial separation would exceed the correlation length. Consequently, it is not necessarily unfeasible to distinguish the 21-cm background, utilizing a radio telescope with a large collecting area. That line radiation allows three-dimensional tomography of the high-z HI renders this a specially interesting technique.

For the 21-cm contribution to be observable, the spin temperature Ts must differ from that of the black-body cosmic background. The gas would be detected in absorption or in emission depending on whether Ts is lower or higher than Trad. The hyperfine levels of HI are affected by the microwave background itself, by collisional processes, and by Lyman α (whose profile is itself controlled by the kinetic temperature). Ts will therefore be a weighted mean of the CMB and gas temperatures.

Before there had been any heat input due to the development of nonlinear structures, the kinetic temperature would be lower than that of the radiation, and the 21 cm would be an absorption feature. In principle, one might be able to detect incipient large-scale structure, even when still in the linear regime, because it leads to variations in the column density of HI, per unit redshift interval, along different lines of sight (5).

When reheating occurs, the situation becomes more complicated (6). The kinetic temperature can rise due to the weak shocking and adiabatic compression that accompanies the emergence of the first (very small scale) nonlinear structure (compare section 2). When photoionization starts, there will also, around each HII domain, be a zone of predominantly neutral hydrogen that has been heated by hard UV or x-ray photons. This latter effect would be more important if the first UV sources emitted radiation with a power-law (rather than just exponential) component.

Because the signal is so weak, there is little prospect of detecting high-z, 21-cm emission unless the signal displays structure on (comoving) scales of several million parsec (Mpc) (corresponding to angular scales of several arc minutes). According to CDM-type models, the gas is likely to have been already ionized, predominantly by numerous ionizing sources, each of subgalactic scale, before such large structures become conspicuous. On the other hand, if the primordial gas were heated by widely spaced, quasar-level sources, each of these would be surrounded by a shell that could feasibly be revealed by 21-cm tomography using, for instance, the new Giant Meter Wave Telescope (7).


2. Models and Methods

To simulate the hazy Archean environment with boundary conditions consistent with recent geochemical constraints, we used a coupled 1-D photochemical-climate model we call Atmos and a 1-D radiative transfer model, SMART (Spectral Mapping Atmospheric Radiative Transfer).

2.1. Coupled photochemical-climate model

Our coupled photochemical-climate model, Atmos, is used to simulate Archean Earth's photochemistry and climate. To use Atmos, the photochemical model (which includes particle microphysics) is run first to generate an initial atmospheric state based on user-specified boundary conditions [gas mixing ratios or fluxes, the solar constant at 2.7 Ga (Claire et al., 2012), the stellar spectral type, total atmospheric pressure, the initial temperature-pressure profile]. Then, the output files from the photochemical model for altitude, pressure, gas mixing ratios, haze particle sizes, and haze number densities are passed into the climate model. The climate model uses the photochemical model's solution as its initial conditions and runs until it reaches a converged state. It then feeds updated temperature and water vapor profiles back into the photochemical model. The models iterate back and forth in this manner until convergence is reached. An example of Atmos finding convergence can be seen in Fig. 1.

FIG. 1. Shown is an example of the Atmos model convergence process. This atmosphere, which has CH4/CO2 = 0.17 and pCO2 = 0.02 (total pressure 1 bar) goes through five coupling iterations. The initial temperature profile it uses was stored from a previous similar atmosphere. Here we show the temperature, water, haze number density, haze particle radii, C2H6 profile, and CH4 profile for each iteration of the coupled model.

2.1.1. Photochemical model

The photochemical portion of the code is based on the 1-D photochemical code developed originally by Kasting et al. (1979), but the version we use here was significantly modernized and updated by Zahnle et al. (2006) and uses the haze formation scheme described by Pavlov et al. (2001b). It was modified by E. Wolf to include fractal hydrocarbon hazes following the methods presented by Wolf and Toon (2010) and was first used to study fractal hazes on Archean Earth by Zerkle et al. (2012). Note that the version of the model used here can simulate atmospheres ranging from extremely anoxic (pO2 = 10 −14 ) to modern-day O2 levels (Zahnle et al., 2006). Subsequent studies using this model or other versions of it to study fractal haze formation include those of Harman et al. (2013), Kurzweil et al. (2013), and Claire et al. (2014), with the latter two of these studies also derived from the same Zahnle et al. (2006) model branch used here. This model also has a long heritage of being used to study photochemistry in nonhazy atmospheres (e.g., Kasting and Donahue, 1980 Pavlov and Kasting, 2002 Ono et al., 2003 Segura et al., 2003, 2005, 2007, 2010 Zahnle et al., 2006 Grenfell et al., 2007 Catling et al., 2010 Domagal-Goldman et al., 2011, 2014 Rugheimer et al., 2013, 2015 Harman et al., 2015 Schwieterman et al., 2016).

The photochemical model parameters are as follows. Our model atmosphere is divided into 200 plane-parallel layers from the surface to 100 km, with a layer spacing of 0.5 km. We show a list of chemical reactions in our Supplementary Table S1 (Supplementary Data are available online at www.liebertonline.com/ast). Our Archean scheme includes 76 chemical species, 11 of which are short-lived (Supplementary Table S2). Short-lived species are considered in photochemical equilibrium (i.e., their atmospheric transport is neglected) and are not part of the Jacobian solved self-consistently at each time step. The mixing ratio of each species is found by solving flux and mass continuity equations in each layer simultaneously using a reverse-Euler method, providing exact solutions at steady state. Vertical transport by molecular and eddy diffusion is included, and boundary conditions that drive the model can be set for each species at the surface and the top of the atmosphere. A δ-2-stream method is used for radiative transfer (Toon et al., 1989). Fixed isoprofiles are assumed for CO2 and N2 in the atmospheres considered here.

Similarly to the work of Zerkle et al. (2012), we set a fixed mixing ratio of CH4 at the surface the model then calculates the surface flux necessary to maintain this mixing ratio. Since haze formation scales with the CH4/CO2 ratio, we find this is the most straightforward way to explore haze thicknesses in our atmospheres. Note that when we discuss CH4/CO2 values in this study, these refer to the ratio at the planetary surface because CH4 does not follow an isoprofile.

Aerosol formation follows the method used in Kasting et al. (1989) and described and updated in Pavlov et al. (2001b). Immediate precursors to haze particles are formed through the reactions C2H + C2H2 → C4H2 + H and C2H + CH2CCH2 → C5H4 + H. Since the full chemical scheme that leads to aerosol formation is not well understood despite both laboratory and theoretical studies (e.g., Hallquist et al., 2009 Hicks et al., 2015), it is assumed that C4H2 and C5H4 condense directly to haze particles (called HCAER and HCAER2 in Supplementary Table S1). In a real atmosphere, the molecules would be larger before aerosols condense, and back-reactions should occur, so this model may overestimate the rate of aerosol formation. Pavlov et al. (2001b) suggested that if the real aerosol formation rate was slower, the atmosphere would compensate by increasing the CH4/CO2 ratio, which would increase the polymerization rate. Further discussion of haze formation pathways and caveats of the approach we use here can be found in Section 4.4. The model's particles form initially with a radius of 0.001 μm. Each layer of the atmosphere has a monomodal size distribution calculated by comparing the coagulation lifetime to the particle removal lifetime via diffusion into another layer or by sedimentation. The aerosols can grow when the coagulation lifetime is longer than the lifetime for removal in a layer.

The maximum radius of a spherical haze particle (i.e., a haze “monomer”) is set to 0.05 μm, the same nominal value used by Wolf and Toon (2010) and similar to the size of the monomers of Titan's fractal haze aggregates (Rannou et al., 1997 Tomasko et al., 2008). Particles larger than this size are treated as fractal agglomerates of nmon spherical monomers of radius Rmon that clump into a larger aggregate with an effective geometric radius Rf given by the relation

Here, α represents a dimensionless constant of order unity, and Df is the “fractal dimension,” which can take on values between 1 and 3. Df = 3 represents a spherical (nonfractal or classical Mie) particle, while Df = 1 represents a string of linearly chained monomers. Titan's fractal aggregates are thought to have a fractal dimension of about 2 on average for the aerosol population (Rannou et al., 1997 Larson et al., 2015). Note that the “effective geometric radius” we refer to above is used only to conceptualize the size of a fractal particle and does not indicate that we use Mie scattering for our fractal particles with the exception of sub-monomer-sized particles (R < 0.05 μm) which remain spherical and thus Mie, we use the mean field approximation for fractal scattering physics for all particles (Botet et al., 1997). The model's fractal production methods are discussed by Zerkle et al. (2012) (including their supplementary online information), where they were first implemented. Additional information about fractal particles and their geometry can be found in the works of, for example, Köylü et al. (1995) and Brasil et al. (1999). The mean field approximation we use for fractal scattering has been validated against scattering by silica fractal aggregates (Botet et al., 1997) and Titan's hazes (Rannou et al., 1997 Larson et al., 2015).

As in the work of Wolf and Toon (2010), the fractal dimension of our particles varies from 1.5 to 2.4 for aggregate particles, and larger aggregates have a larger fractal dimension to account for folding as the particles coagulate. In general, compared to spherical particles, fractal particles produce more extinction in the ultraviolet (UV) but less in the visible and near infrared (NIR). In addition, fractals tend to be more forward scattering in the visible and NIR and more isotropically scattering in the UV compared to equal-mass spherical particles. Their weakened visible extinction and enhanced forward scattering compared to spherical particles means they produce less cooling since they scatter less incident sunlight back to space (see Fig. 3 in Wolf and Toon, 2010). Figure 2 shows the extinction efficiency (Qext) and single-scattering albedo of different fractal particle sizes together with the haze optical constants we adopt in this study (Khare et al., 1984a). A discussion of our choice of optical constants and comparison to others in the literature can be found in Section 4.5.

FIG. 2. The top panels present the extinction efficiency (Qext) and single-scattering albedo ( = Qscat/Qext) of four sizes of fractal hydrocarbon particles used in this study and in Wolf and Toon (2010). The spherical monomers comprising these particles are 0.05 μm in radius. The radii on the plot correspond to the radii of equivalent-mass spherical particles, and the fractal dimensions of these particles, from smallest to largest, are 3 (spherical), 1.51, 2.28, and 2.40. The number of monomers in these particles are 1, 8, 1000, and 8000. These particles tend to scatter and absorb light more efficiently at shorter wavelengths, and larger particles have flatter wavelength dependence for the scattering efficiency. Refractive indices, shown in the bottom panels, are presented from information in Khare et al. (1984a).

FIG. 3. The gas profiles for H2O, CH4, CO, CO2, and C2H6 for planets with pCO2 = 0.01 bar for CH4/CO2 = 0.1 (on the left) and CH4/CO2 = 0.2 (on the right). Also shown are the profiles for the haze particle number density (in pale orange). The CH4/CO2 = 0.1 haze profile is divided by 1000, and the CH4/CO2 = 0.2 haze profile is divided by 1 × 10 5 in order to plot it on the same axis as the gases. The profiles in the right panel show larger amounts of CH4, H2O, and C2H6 above 60 km in altitude and illustrate how haze-induced shielding can prevent photolysis of these gases. The sharp decrease in haze particle number density between 60 and 70 km in the right panel shows where fractal coagulation occurs. The atmosphere above the fractal coagulation region is populated by spherical submonomer particles.

In the version of the photochemical model used here, we corrected an error relating to the calculation of the number of C5H4 molecules composing HCAER2 haze particles. Previously, the model calculated the number of molecules per HCAER2 particle inappropriately using the mass of C4H2 instead of C5H4. In addition, we added more particle sizes to the model's scattering grid, increasing the number from 34 particle sizes to 51, and we added options to use different monomer sizes and optical constants than the ones used here for our nominal haze study how variation of these parameters impacts haze formation is a subject of future work. Gas mixing ratios at the surface can be more finely tuned than in previous versions of the model from the addition of a significant figure to the species boundary conditions input file.

The photochemical model is considered converged when redox is conserved and a re-run of the model using last run's output as initial conditions occurs quickly (i.e., <50 time steps).

2.1.2. Climate model

Our climate model was originally developed by Kasting and Ackerman (1986). The model we use here has evolved considerably since its first incarnation, and versions of it have been applied in subsequent studies on varied topics such as the habitable zones for several stellar spectral types (Kopparapu et al., 2013), the climate of early Mars (Ramirez et al., 2013), the atmospheres of Earth-like planets around various stellar types (Segura et al., 2003, 2005, 2010 Rugheimer et al., 2013), clouds in exoplanet atmospheres (Kitzmann et al., 2010, 2011a), and the climate of early Earth (Haqq-Misra et al., 2008). The version we use here is based directly on that used by Kopparapu et al. (2013). It uses a correlated-k method to compute absorption by spectrally active gases (O3, CO2, H2O, O2, CH4, and C2H6). This model has CO2 and H2O correlated-k coefficients updated as described by Kopparapu et al. (2013). Our older CH4 coefficients may overestimate the surface temperature by ≲5 K at the CH4 mixing ratios used here (Byrne and Goldblatt, 2015). However, as we discuss in Section 4.2, our model underpredicts the Archean temperature by about 2–5 K compared to 3-D climate models with more complete physics describing the planetary system, so these two effects may cancel each other out. The aforementioned gas profiles are passed to the climate model from the photochemical model when running in coupled mode. The net absorbed solar radiation in each layer of the atmosphere is computed using a δ-2-stream multiple scattering algorithm (Toon et al., 1989) spanning from λ = 0.2 to 4.5 μm in 38 spectral intervals. For net outgoing IR radiation, we use a separate set of correlated-k coefficients for each gas in 55 spectral intervals spanning wave numbers of 0–15,000 cm −1 .

We have made several modifications to the climate model used here. The model previously incorporated the spectral effects of spherical hydrocarbon particles, and it has been updated in our study to include fractal hydrocarbon scattering efficiencies using the mean field approximation of Botet et al. (1997) discussed previously. We have also updated the model so that haze profiles can be passed to it from an input file or by the photochemical code in previous versions of the climate model, haze distributions were hard-coded and had to be edited manually. We corrected a discrepancy in the spacing between atmospheric layers in the routine that outputs coupling files for the photochemical model: our photochemical model layer spacing is 0.5 km, but a layer spacing of 1 km had been hard-coded. Coupling subroutines have been improved to be able to accept information about atmospheric pressure, stellar parameters, and haze parameters as input from the photochemical model. We also added options to turn ethane opacity and 1-D ice-albedo feedbacks (described in Section 4.1.1) on or off.

We have been unable to run the climate model to convergence using the same top-of-atmosphere pressure used for the photochemical model: the photochemical model extends to 100 km, but we have only been able to successfully run the climate model up to about 80 km for our 1 bar atmospheres. Thus, when temperature and water profiles are passed from the climate model to the photochemical model, they become isoprofiles above the top of the climate grid based on the highest-altitude temperature from the climate grid calculations. At these altitudes the atmosphere is thin, and the particles are very small both of these effects lead to this portion of the atmosphere having little impact on radiative transfer and climate. We performed a sensitivity test of how the temperature at these altitudes affects the resultant haze distribution in the photochemical model, and the sizes of the largest haze particles produced by an atmosphere that becomes an 80 K isotherm above 80 km versus a 150 K isotherm differ by less than 5%. In the climate model, shifting the particles in Fig. 1 above 80 km down to lower altitudes alters the surface temperature by <0.5 K.

The climate model is considered converged when the change in temperature between time steps and change in flux out the top of the atmosphere are sufficiently small (typically on the order of 1 × 10 −5 ).

2.2. The SMART model

To generate synthetic spectra for the atmospheres we produce with Atmos, we feed outputs from the Atmos model (the temperature-pressure profile, gas mixing ratio profiles, and the haze particle profile), into the SMART code, a 1-D line-by-line fully multiple scattering radiative transfer model (Meadows and Crisp, 1996 Crisp, 1997). SMART has been validated against observations of multiple solar system planets (Robinson et al., 2011 Arney et al., 2014). The Line-by-Line Absorption Coefficients (LBLABC) code, a companion to SMART, creates line-by-line absorption files for input gas mixing ratios and temperature-pressure profiles using HITRAN 2012 line lists (Rothman et al., 2013). SMART can also incorporate aerosols: as input, it requires “cloud files” with altitude-dependent opacities as well as the particle asymmetry parameter and the extinction, scattering, and absorption efficiencies (Qext, Qscat, and Qabs). For spherical particles (our small monomers), we use the code “Miescat,” to calculate these efficiencies using the indices of refraction measured by Khare et al. (1984a). For fractal hydrocarbon particles, we use scattering inputs from the Wolf and Toon (2010) photochemical study generated with the fractal mean field approximation (Botet et al., 1997). Spherical particles use a full Mie phase function, while fractal particles employ a Henyey-Greenstein phase function (Henyey and Greenstein, 1941). To generate transit transmission spectra, we use the SMART-T model (Misra et al., 2014a, 2014b). This version of SMART uses the same inputs as the standard code but simulates the longer path lengths and refraction effects associated with transit transmission observations.

To create SMART cloud files from Atmos haze outputs, we have written a script that bins the haze particles generated by the photochemical model into specified radii (also called particle “modes”) while preserving the total mass of each atmospheric layer. The particle mode sizes we use span from 0.001 to 2 μm larger particles do not exist in our atmospheres due to rainout. Spherical modes are R = 0.001, 0.005, 0.01, and 0.05 μm. Fractal modes are R = 0.06–2 μm with four modes between 0.06 and 0.1, 10 equally spaced modes between 0.1 and 1 μm, and 2 μm. In total, this represents 19 particle modes.

In each layer of the SMART cloud files, we include a mixture of two particle modes the mass density contributed by the two modes is selected based on the distance in log space of the Atmos particle radius to each neighboring SMART size bin. For example, if Atmos produces a particle of radius 0.33 μm in a layer, the corresponding layer in SMART will include 0.3 and 0.4 μm particles each comprising 50% of the layer's mass. This binning is necessary because the photochemical model generates many dozens of finely differentiated haze particle radii, but SMART model run time with this many particle sizes is infeasible.

Once we have binned the Atmos particle radii to our SMART size grid, we must compute the total optical depth from each particle mode at a reference wavelength in each atmospheric layer. We arbitrarily select 1 μm as our reference wavelength. Optical depth in a layer, τ, from particles of a given radius, R, depends on the number density of particles per particle size, n(R), the thickness of the atmospheric layer, z, and the wavelength-dependent extinction efficiency, Qext:

For fractal particles (R > 0.05 μm), the cross-sectional area and the corresponding extinction efficiencies are computed relative to the radius of an equal-mass spherical particle, following the conventions of mean-field approximation (Botet et al., 1997). Spherical particles in SMART are binned according to log-normal size distributions using the radii mentioned previously and a mode standard deviation of 1.5, which is realistic for an aerosol distribution (Tolfo, 1977). For fractal particles, we use a monodisperse distribution, the same size distribution used to compute our inputs from the previous Wolf and Toon (2010) fractal haze study and the same distribution used in the Atmos model.

2.3. Model inputs

In the photochemical model, we set a haze monomer density of 0.64 g/cm 3 , which is consistent with the laboratory results of Trainer et al. (2006) for early Earth. This density is used in the model to calculate the masses of haze particles and is updated from the value of 1 g/cm 3 used by previous studies employing our photochemical model. Hörst and Tolbert (2013) measured a similar effective particle density, 0.65 g/cm 3 , for a 0.1% CH4 haze experiment using a UV lamp. One-tenth percent CH4 is consistent with the atmospheres we simulate, although the Hörst and Tolbert hazes were Titan analog simulants lacking the CO2 present in the Trainer et al. experiments. We apply a Manabe/Wetherald relative humidity model for the troposphere (Manabe and Wetherald, 1967) with a surface relative humidity of 0.8 in both the climate and photochemistry models. This humidity parameterization is further described by Pavlov et al. (2000). Our Archean simulations use the solar constant at 2.7 Ga (0.81 = S/S0, where S0 is the modern solar constant and S is the solar constant at 2.7 Ga) modified by a wavelength-dependent solar evolution correction (Claire et al., 2012). We chose this time because it corresponds to the age of the constraints on CO2 used by our study (Driese et al., 2011). We set the mixing ratio of O2 at the surface to 1.0 × 10 −8 , consistent with the Zerkle et al. (2012) study. These conditions reflect the time period after the evolution of oxygenic photosynthesis but prior to Earth's GOE in which substantial biogenic fluxes of both oxygen and methane would have vented into a predominantly reducing atmosphere (Claire et al., 2014). Unless otherwise specified, the surface albedo used by the climate model is 0.32. This includes the effect of clouds, which is standard in this 1-D treatment (Kopparapu et al., 2013) and is the albedo that reproduces the average temperature of present-day Earth (288 K) with modern atmospheric conditions. Of course, the true cloud distribution on Archean Earth is unknown, and clouds may have had important climatic effects on our early planet (Goldblatt and Zahnle, 2011). The solar zenith angles (SZAs) used in the climate and photochemical models were chosen to best represent globally averaged behavior of the physics in each specific model, which Segura et al. (2003) found as SZA = 45° in the photochemical model and SZA = 60° in the climate model. These zenith angles are both tuned to reproduce modern-day Earth's average chemical profiles and climate, respectively.

For our SMART spectral simulations, our nominal spectra assume an ocean surface albedo (McLinden et al., 1997). In cases where an icy surface is used, we use an albedo from the USGS Digital Spectral Library (Clark et al., 2007). Our solar spectrum was modeled by Chance and Kurucz (2010) and was scaled by the solar evolution model (Claire et al., 2012) mentioned previously. The SZA is set at 60° for the reflection spectra, which approximates a planetary disk average near quadrature (planet half illuminated to the observer).


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Abstract

The onset and nature of the earliest geomagnetic field is important for understanding the evolution of the core, atmosphere and life on Earth. A record of the early geodynamo is preserved in ancient silicate crystals containing minute magnetic inclusions. These data indicate the presence of a geodynamo during the Paleoarchean, between 3.4 and 3.45 billion years ago. While the magnetic field sheltered Earth’s atmosphere from erosion at this time, standoff of the solar wind was greatly reduced, and similar to that during modern extreme solar storms. These conditions suggest that intense radiation from the young Sun may have modified the atmosphere of the young Earth by promoting loss of volatiles, including water. Such effects would have been more pronounced if the field were absent or very weak prior to 3.45 billion years ago, as suggested by some models of lower mantle evolution. The frontier is thus trying to obtain geomagnetic field records that are ≫ 3.45 billion-years-old, as well as constraining solar wind pressure for these times. In this review we suggest pathways for constraining these parameters and the attendant history of Earth’s deep interior, hydrosphere and atmosphere. In particular, we discuss new estimates for solar wind pressure for the first 700 million years of Earth history, the competing effects of magnetic shielding versus solar ion collection, and bounds on the detection level of a geodynamo imposed by the presence of external fields. We also discuss the prospects for constraining Hadean–Paleoarchean magnetic field strength using paleointensity analyses of zircons.


Line-of-Sight Anisotropies in the Cosmic Dawn and Epoch of Reionization 21-cm Power Spectrum

The line-of-sight direction in the redshifted 21-cm signal coming from the cosmic dawn and the epoch of reionization is quite unique in many ways compared to any other cosmological signal. Different unique effects, such as the evolution history of the signal, non-linear peculiar velocities of the matter etc. will imprint their signature along the line-of-sight axis of the observed signal. One of the major goals of the future SKA-LOW radio interferometer is to observe the cosmic dawn and the epoch of reionization through this 21-cm signal. It is thus important to understand how these various effects affect the signal for its actual detection and proper interpretation. For more than one and half decades, various groups in India have been actively trying to understand and quantify the different line-of-sight effects that are present in this signal through analytical models and simulations. In many ways the importance of this sub-field under 21-cm cosmology have been identified, highlighted and pushed forward by the Indian community. In this article, we briefly describe their contribution and implication of these effects in the context of the future surveys of the cosmic dawn and the epoch of reionization that will be conducted by the SKA-LOW.

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3 Temporal Evolution of ‘Oumuamua

We demonstrated above that a pure N2 ice composition is capable of reproducing the observed non-gravitational acceleration at a distance of around 1.4 au from the Sun for an albedo of 0.64. At that time, however, ‘Oumuamua would have had a mass of only 8.01 × 10 6 kg while losing mass at a rate of 0.37 kg/s (3.2 × 10 4 kg/day) such that even between 1.4 and 2 au it would have shrunk by around 0.7 m. It is thus clear that ‘Oumuamua would have evolved substantially over time.

3.1 Passage through the Solar System

As we can see from Equation 5, the rate of mass loss from ‘Oumuamua is a strong function of the heliocentric distance. From our starting point at 1.42 au we can integrate forwards and backwards to find the evolution of the size and shape of ‘Oumuamua over time. In Figure 2 we show the evolution in the mass, surface temperature, and axis ratio (a/c) of ‘Oumuamua over a period of around 18 months before and after perihelion alongside its distance from the Sun for comparison. Times are measured relative to October 27, 2017 at which epoch we fix the distance from the Sun as 1.42 au and the size as 45.5 m × 43.9 m × 7.5 m. The albedo is set at 0.64, which we assume remains constant. For the orbital evolution we assume a semi-major axis of −1.2978 au and eccentricity 1.19951 for which perihelion passage occurs 48 days before our fixed point (September 9, 2017) (From JPL HORIZONS service https://ssd.jpl.nasa.gov/sbdb.cgi?sstr=2017U1). The non-gravitational acceleration is never large enough to modify the orbit sufficiently to significantly alter our calculations, and so for the purposes of Figure 2 we neglect the non-gravitational acceleration. We assume that in a time interval Δt each semi-axis of the triaxial ellipsoid decreases by an amount h = (dR/dtt, and the surface area of the ellipsoid is re-computed at the end of each timestep.

Change in ‘Oumuamua’s parameters over time. Clockwise from top left: distance from the Sun, mass, axis ratio (a/c), surface temperature. Times are relative to October 27, 2017. Perihelion occurs at −48 days.

Unsurprisingly, the period immediately around perihelion dominates the change in mass and axis ratio. In the 50 days either side of perihelion passage the mass drops by a factor of 10 while the axis ratio rises from just above 2:1 to 6:1. It is notable, however, that the evaporative cooling due to the extreme mass loss is so effective that the surface never rises above 47 K. Outside this narrow window, mass loss continues at a much lower rate, and the axis ratio continues to undergo significant evolution because the c axis has shrunk to such a small size. By the time ‘Oumuamua passed the orbit of Uranus in September 2020 it would have dropped in mass by roughly another factor of 2, to 4.81 × 10 6 kg, and reached an axis ratio of 8.3:1.

As the mass and mass loss changed, so too would have the non-gravitational acceleration. Micheli et al. ( 2018 ) found that the observed non-gravitational acceleration obeyed a relationship that lies between d −1 and d −2 , probably closer to d −2 . We can immediately see from Equation 6 that the insolation induces a d −2 dependence, with a non-linear deviation for the thermal emission term. In Figure 3 we show the change in our predicted non-gravitational acceleration as a function of time during a window of a few months spanning the range of observations that were used by Micheli et al. ( 2018 ) to fit the non-gravitational acceleration, these observations were densest from October 18 to November 23, 2017 with precovery data for October 14, 2017, and additional points on December 12, 2017 and January 2, 2018. Not only does our predicted acceleration match the magnitude of the non-gravitational acceleration at 1.42 au, but over the relevant range our predicted curve is very close to d −1.8 , which we consider a good fit to the observations given the uncertainties involved, both in the observations and in some of the parameters in our equations (e.g., what value of τ is appropriate, and corrections for the ellipsoidal shape of ‘Oumuamua).

Non-gravitational acceleration predicted by our model (solid red) for time around October 27, 2017 as compared with different relationships that are constant power laws in d. At left we show the predicted non-gravitational acceleration while at right we multiply by (d/au) 2 to provide a more detailed view of the differences between the curves. The dashed and dotted black lines show d −2 and d −1 relations respectively, while the solid black line shows a d −1.8 relationship, which provides a good fit to our prediction over the relevant range. The range of observations over which the acceleration was fit by Micheli et al. ( 2018 ) runs from -13 to +68 days (14 October 2017 to 2 January 2018).

Our scenario is also consistent with ‘Oumuamua’s rotation. ‘Oumuamua is tumbling, in non-principal axis rotation Fraser et al. ( 2018 ), but spinning only once per 8 h or so. This “slow” rotation rate of ‘Oumuamua has been taken by Rafikov ( 2018b ) as evidence against outgassing, arguing that the forces needed to provide the non-gravitational acceleration would torque and spin up the object until it underwent rotational fission. The underlying picture is one like a normal comet, in which most of the jetting occurs from isolated spots on the surface, providing a torque that increases the rotation rate Ω at a rate dΩ/dt = ζ(RM/I)a, where a is the non-gravitational acceleration, R, M and I the characteristic size, the mass, and the moment of inertia of the object, ζ a dimensionless number such that ζR is the effective lever arm.

Based on his analysis of seven regular comets, Rafikov ( 2018a ) derived an average value of log ζ = −2.21 ± 0.54. Mashchenko ( 2019 ) found a best fit of ‘Oumuamua’s light curve if it were oblate and experienced a torque consistent with log ζ = −2.34, which, as they pointed out, is within the range of values Rafikov ( 2018a ) inferred for comets. Since IMR 2 , the spin-up rate is proportional to R −1 and therefore p 1/2 . For a given spin-up rate, ζ is proportional to p −1/2 . Because we are arguing for an albedo p = 0.64 rather than the value p = 0.1 assumed by Mashchenko ( 2019 ) when deriving ζ, we favor a smaller body more easily spun up, and therefore a smaller value of the effective lever arm coefficient, log ζ ≈ − 2.74. However, this is still within the range of values observed among comets (Rafikov 2018a ). Moreover, a lower value of ζ is consistent with the idea of ‘Oumuamua as a monolith of pure N2 ice without localized jetting, in accordance with our assumption of sublimation across the hemisphere, and our use of the coefficient 1/3, in our derivation of the jetting force. We note that Seligman et al. ( 2019 ) also examined the possible spin-up of ‘Oumuamua and found that it is possible for the rotation rate to oscillate around the observed ∼8 h value under the action of sublimation jetting for appropriate venting angles, however it is not clear how their results scale to a body with a substantially different albedo. It has been argued that for ‘Oumuamua to aquire its tumbling motion would require many Gyr of passage through the interstellar medium (Zhou, 2020 ), but it seems clear that ‘Oumuamua must have experienced significant (albeit lower than is typical for comets) torques within the Solar System, as it lost ∼92% of its mass.

3.2 Passage through the Interstellar Medium

An important factor in ruling out H2 as a likely composition for ‘Oumuamua is that an H2 ice body would experience rapid erosion during its passage through the interstellar medium (ISM). Hoang and Loeb ( 2020 ) discussed this erosion in detail, showing that even a multi-km H2 ice body would be completely eroded away in less than 10 8 years. For H2, Hoang and Loeb ( 2020 ) found that simple thermal sublimation was dominant since the equilibrium surface temperature of a body in the ISM is barely below the sublimation temperature of H2 ice. The sublimation temperature of N2 ice is about a factor of 7 higher than that of H2 ice and the exponential dependence of sublimation (Equation 2) makes it immediately apparent that direct, thermally driven sublimation will be negligible at ISM temperatures for N2 but it is nonetheless prudent to consider other possible erosion mechanisms.

Domokos et al. ( 2009 ) described how isotropic abrasion by dust grains impacting the surface of a body can lead to an increase in the axis ratio, much as we have described above for outgassing, and Domokos et al. ( 2017 ) attributed ‘Oumuamua’s shape to abrasion by dust grains eroding its surface as it passed through the ISM. However, ‘Oumuamua is unlikely to encounter sufficient material as it passes through the ISM to result in significant change to its mass and dimensions. For a typical ISM density of around 1 proton per cm 3 and a dust-to-gas mass ratio of 0.01, a body with mean diameter ∼50 m and a relative velocity ∼10 km/s will only collide with around 10 3 kg per Gyr of matter in total, and only around 10 kg per Gyr of dust, a tiny fraction of the ∼ 10 7 kg mass of the body. Even traveling 10 pc through a giant molecular cloud with a mean density of 10 3 protons per cm 3 would only result in encountering around 10 3 kg of gas and dust. Dust abrasion is thus clearly insufficient to result in any change to ‘Oumuamua’s size or shape.

Another possible mechanism is photodesorption. While both visible and UV photons have sufficient energy to overcome the desorption energy of an N2 molecule in N2 ice (∼0.07 eV) the efficiency of the process is only around 5 × 10 −3 N2 molecules per photon, since N2 ice is an inefficient absorber at most optical and UV wavelengths (Bertin et al., 2013 Fayolle et al., 2013 ). We assume the interstellar radiation field will deliver photons with an energy flux of 2.7 × 10 −6 W m −2 , each with a typical energy of around 10 eV (Mathis et al., 1983 ). This would produce a desorption rate of 8 × 10 9 N2 molecules m −2 s −1 , or an erosion rate of just over 1 cm/Gyr. As with dust abrasion, this is insufficient to result in any substantial changes in ‘Oumuamua’s size.

Finally we consider galactic cosmic rays (GCRs). Integrating over energy, the GCR proton and alpha-particle energy flux in the ISM near the Sun is around 1.9 × 10 −5 W m −2 (using the analytical formula of Webber ( 1998 ), but scaled down by a factor of 1.5 to better match the observations compiled by Tatischeff et al. ( 2014 )). The majority of the incident particles would have energies in the range 10–100 MeV/nucleon. Vasconcelos et al. ( 2017 ) measured the erosion of an N2:CH4 ice (95:5 mass ratio) by 15.7 MeV oxygen ions (1 MeV/nucleon) and found that after receiving a fluence of 6 × 10 17 ions m −2 the ice was reduced in thickness by about 8 μm. The ions deposited 930 keV/μm, or a total of 7.4 MeV each. Based on this experimental data we infer the removal of 1 N2 molecule for roughly every 26 eV delivered by GCRs. We note that the stopping lengths of typical GCRs will not exceed a fraction of a cm and so assume there is no reduction in erosion efficiency for higher energy GCRs. For an interstellar GCR energy flux of 1.9 × 10 −5 W m −2 we calculate an average erosion rate of 6.6 m/Gyr. For N2 ice, GCR erosion is thus the dominant erosion mechanism in the ISM and can potential alter the size and shape of the body over long periods. This erosion rate due to GCRs is about the same as we calculate for the thermal sublimation due to solar radiation at a distance of about 130 au, so beyond around 130 au from the Sun GCR erosion dominates over thermal sublimation. This is roughly the same distance as the heliopause, within which GCRs are suppressed by the Solar magnetic field (Gurnett et al., 2013 ).

It is important to note that the GCR erosion rate of 6.6 m/Gyr corresponds to the GCR flux in the neighborhood of the Sun, today. The GCR flux is roughly proportional to the star formation rate, and so the GCR flux can be expected to have tracked variations in the star formation rate in the vicinity of the Sun over time. The Sun is currently in an inter-arm region with a GCR flux characteristic of much of the Galaxy, but the star formation rate in the spiral arms is substantially higher, such that the GCR flux is expected to be around 3.7 times higher within a spiral arm than the galactic average (Dunham et al., 2020 Fujimoto et al., 2020 ). Over a period of a Gyr, traveling at 9 km/s relative to the local standard of rest, ‘Oumuamua could travel tens of kpc, likely passing in and out of spiral arms multiple times. The analysis of Vallée ( 2005 ) suggests that the Sun spends roughly 50% of its time within 1 kpc of spiral arms (the diffusion length of GCRs), and the other 50% in the inter-arm regions. This implies an average erosion rate over the last few hundred Myr that is around a factor of 2.4 times higher than the rate in the Solar neighborhood today, that is ∼15.4 m/Gyr. In addition, the star formation rate was higher in the past across the entire Galaxy: it was greater than 5 times the present rate over 8 Gyr ago, falling to about twice the current rate 5–6 Gyr ago, peaking again at around 5 times the present value 2–3 Gyr ago, and finally falling since then to the present value (Mor et al., 2019 ). Over the 4.5 Gyr since the birth of the Solar system the total erosion would have been 260 m along each semi-axis, averaging 57 m/Gyr. In the last 2 Gyr ‘Oumuamua would have eroded by 92 m at an average of around 3 times the present rate, and even over just the last Gyr it would have eroded at an average of about twice the present rate (including the correction for spiral arms) for a total of around 31 m.

How long ‘Oumuamua has been traveling through the ISM is not known, but Almeida-Fernandes and Rocha-Pinto ( 2018 ) placed an upper limit of around 1.9–2.1 Gyr based on its low velocity dispersion relative to the LSR. Had ‘Oumuamua been in interstellar space for this maximum time of around 2 Gyr, it would have been eroded by over 90 m in radius, implying an initial mass upon ejection from its stellar system of around 8 × 10 9 kg. This would mean that ‘Oumuamua entered the solar system with only 1% of the mass it had when it left its parent system while this is not impossible, it seems unlikely. By comparison, if ‘Oumuamua departed its parent system around 0.4–0.5 Gyr ago, it would have been eroded by around 10 m along each semi-axis, entering the Solar system with slightly under half of its initial mass, a much more plausible value. Traveling at 9 km/s for around 0.4–0.5 Gyr, ‘Oumuamua could have traveled about 4 kpc, albeit not in a straight line: its motion through the Galactic potential would have changed its velocity en route, and this distance must include epicyclic motions. Since a young stellar system is the most likely candidate to be ejecting large quantities of material we tentatively suggest an origin around 0.4–0.5 Gyr ago in the Perseus spiral arm, which is about 2–3 kpc from the Sun (Kounkel et al., 2020 ) and consistent with ‘Oumuamua’s approach from the direction of Vega. Using this starting point we provide a summary of the mass and dimensions of ‘Oumuamua at various epochs along its journey from its parent system, to, and through the Solar system in Table 1. We note that if ‘Oumuamua were traveling through the ISM for 0.4–0.5 Gyr, it would have seen its axis ratios increase from a/c = 1.7 to 2.1. Typical axis ratios of fragments in the solar system are a: b: c ∼ 2 : 1.4: 1, never exceeding 3:1 (Domokos et al., 2017 ). While it would not be implausible for ‘Oumuamua to have been ejected from its stellar system with an axis ratio of 2.1, an initial axis ratio 1.7:1, consistent with travel through the ISM for 0.4–0.5 Gyr, is apparently more likely.

Epoch Distance from sun (au) Time Tsurf (K) 2a (m) 2b (m) 2c (m) a/c Mass (× 10 6 kg)
Ejection from parent system - ∼0.45 Gyr ago - 92.4 90.8 54.4 1.70 244
130 14 Apr. 1995 19.9 72.4 70.8 34.5 2.10 94.4
30 1 Dec. 2012 30.1 72.4 70.8 34.4 2.10 94.2
5.2 16 Jan. 2017 35.0 71.7 70.1 33.8 2.12 90.6
Perihelion passage 0.255 9 Sep. 2017 47.1 57.6 56.1 19.7 2.92 34.0
Optical observations 1.42 27 Oct. 2017 39.4 45.4 43.9 7.50 6.06 7.98
Spitzer observations 2.0 21 Nov. 2017 38.1 44.7 43.2 6.81 6.57 7.03
5.2 2 May 2018 35.0 43.6 42.0 5.64 7.72 5.52
30 18 June 2022 30.1 42.9 41.3 4.96 8.66 4.65
130 14 Feb. 2040 19.7 42.8 41.3 4.92 8.72 4.64

Program

Welcome and opening remarks on the status of the Chandra X-ray Observatory.

Status of Lynx

X-rays and the galaxy cluster landscape for astrophysics and cosmology

Studies of galaxy clusters have proved crucial in helping to establish the standard model of cosmology, and X-ray observations have been central to this work. I will summarize the latest results on cosmology from cluster studies, highlighting Chandra's key role in providing precise, relative mass calibration. The prospects for progress over the next decade and beyond are outstanding with new cluster catalogs, hundreds of times larger and with far greater redshift reach, being constructed across a variety of wavelengths. X-ray follow-up observations will be vital to the full exploitation of these catalogs. I will discuss ways in which to make these contributions as efficient and impactful as possible.

Probing the Detailed Physics of Galaxy Cluster Plasmas with Lynx: Predictions from Mock Observations

The intracluster medium (ICM) of clusters of galaxies is the largest example of a space plasma. The high angular resolution of the Chandra X-ray Observatory has revealed a wealth of complex structure in the ICM, which in principle can be used to probe its plasma physics and kinematics. The microphysical properties of the plasma, such as its viscosity and thermal conduction, are still poorly constrained. It is known from radio observations that the ICM is magnetized, but no evidence has yet been seen of dynamical effects of the weak magnetic field on the gas, despite the fact that simulations predict that the magnetic field should grow strong enough in some cases to have an effect. Observations of surface brightness fluctuations and cold fronts with Chandra have also provided indirect evidence of complex gas motions, including bulk flows and turbulence. The first direct measurements of such gas motions in the Perseus cluster were made by Hitomi with its microcalorimeter, but its low angular resolution severely limits the conclusions that can be drawn about the nature of these motions. In order to provide answers to these questions, a mission with comparable angular resolution to Chandra but higher effective area and higher spectral resolution is needed. I will present results from mock observations of MHD simulations of the ICM to show how Lynx will constrain the detailed physical and kinematical properties of the cluster plasma. In particular, Lynx will have the capability to detect the effects of strong magnetic fields on the cluster plasma, definitively resolve hydrodynamical instabilities to place limits on the ICM viscosity, and map the velocity structure of the gas down to small scales.

Hydrodynamical Cosmological Simulations of Galaxy Cluster Outskirts

In recent years, the outskirts of galaxy clusters have emerged as a frontier to study the evolution of galaxy clusters and the intergalactic medium. In this talk, we will present results from a mass-limited sample of 65 galaxy clusters from the Omega500 cosmological hydrodynamical simulations and corresponding mock Chandra/Lynx X-ray observations. With these results, we will discuss the prospects of probing physical processes in cluster outskirts (such as gas motions, gas clumping, and ion equilibration) with recent deep Chandra XVP observations of A133, large samples of SZ-selected clusters, and the proposed Lynx mission.

From the Chandra Deep Fields to Lynx

As the deepest X-ray surveys ever performed, the Chandra Deep Fields have provided fundamental insights about active galactic nuclei (AGNs) and starburst/normal galaxies in the distant universe. Together they have detected more than 1800 faint X-ray sources that can be characterized in detail using the unmatched deep multiwavelength coverage in these sky regions e.g., about 98% of the sources have multiwavelength counterparts and spectroscopic/photometric redshifts. I will briefly summarize some of the main recent results from the Chandra Deep Fields, focusing on the most-distant AGNs at z&sim3-5 and available constraints upon the seeds of the first supermassive black holes (SMBHs) in the Universe. I will then describe how these studies can be advanced in the relatively near-term, combining Chandra and XMM-Newton data with observations by, e.g., ALMA, JWST, LSST, and new instruments on large ground-based telescopes. I will end by highlighting the advances expected from Athena and the unique, critical role that Lynx surveys will play by detecting large numbers of SMBH seeds down to &sim30,000 solar masses at z&sim10.

Resolving the hidden connections between black holes, galaxies, and halos with Chandra and Lynx

Despite remarkable recent progress, connecting the growth of black holes to the evolution of their host galaxies and dark matter halos has remained challenging. AGN exhibit complex variability in luminosity, accretion processes, and obscuration over a wide range of timescales, so we require large X-ray and multiwavelength surveys to obtain a complete picture via statistical studies of AGN host galaxies and spatial clustering. I will present recent studies by our group and colleagues that use extragalactic surveys (particularly with Chandra, NuSTAR, and WISE) to uncover as much as possible of the complete population of AGN and study their hosts and large-scale structures. These results suggest that AGN accretion is a surprisingly universal, yet highly stochastic process, and uncover an underlying connection between the gas supply that fuels star formation and the accretion rates and obscuration of the growing black holes. I will outline the exciting science in this area that will be enabled by the upcoming Chandra Deep Wide Field Survey, as well as forecasts for the future with NASA's Lynx concept X-ray mission. This work is supported in part by NASA through grant numbers NNX15AP24G and NNX15AU32H, and the National Science Foundation through grant number 1554584.

A 5 per cent determination of the local black hole occupation fraction

I will discuss the feasibility of a few per cent level measurement of the local black hole occupation fraction through Lynx imaging observations of local volume galaxies. I will provide quantitative estimates of how the occupation fraction accuracy and observational strategy vary as a function of a series of parameters, including the HDXI angular resolution, exposure time, distance, and host galaxy stellar mass. Most notably, I will show that 0.1 arcsec resolution will yield a 1-5 per cent accuracy in the occupation fraction below stellar masses of 1e+10 Msun with observations of about 3,500 galaxies within 100 Mpc. This measurement can be efficiently carried out by combining dedicated, snapshot observations of low mass galaxies with a commensal survey approach. It will constitute a long-lasting legacy for the mission by establishing a benchmark for any model which aims to reproduce the assembly of galaxies and their nuclear black holes. Concurrently, it will yield an independent constraint to black hole seed formation models, complementing orthogonal efforts which will be carried out at high red-shifts.

Identifying First X-ray Sources

Detection of X-ray emission from the first population of sources could constrain formation of super-massive black holes and properties of the first population of X-ray binaries. However, direct observations of high redshift populations require large integration times and are highly biased because only the brightest objects at high redshifts can be detected. A useful measurement is of the unresolved cosmic X-ray background (CXB) which constrains both the luminosity of the most common objects and clustering of these sources. An alternative way to probe the first population of X-ray sources is via its effect on the environment. These sources had a dramatic effect on the Universe heating and mildly ionizing the intergalactic medium. One of the most efficient tools to probe the thermal state of the IGM at high redshifts is the radio signal of neutral hydrogen with the rest-frame wavelength of 21 cm. In my talk I will discuss how cross-correlation between cosmic X-ray background and the 21 cm signal can be used to constrain the first population of X-ray sources.

New Measurement of the joint unresolved CIB vs CXB fluctuations: a preview of Lynx sources.

We present a new study of the source subtracted Spitzer-IRAC Cosmic Infrared Background (CIB) and Chandra-ACIS Cosmic X-ray Background (CXB) surface brightness fluctuations cross-power spectra. Our investigation used data from the Chandra Deep Field South (CDFS), Hubble Deep Field North (HDFN), EGS/AEGIS field and UDS/SXDF surveys for a total of 1160 Spitzer hours and &sim12 Ms Chandra data over a total area of 0.3 deg 2 . For the first time we are able to detect signal on the angular scales >20 &Prime between the 3.6&mum, 4.5&mum and [0.5-2] keV bands at &sim5&sigma and 6.3&sigma significance, respectively. We reconstructed the contribution of known unmasked source population at z<6 and found that, with respect to it, our signal is excess of about an order of magnitude at 5&sigma level accounting for &lesssim5% of the CXB flux. The level of coherence between the two background fluctuation is &Cscr&sim0.2. We discuss a possible interpretations of such excess in term of direct collapse black holes, primordial black holes, scattering in the interstellar medium and intra halo light.

Missing Baryons Around Galaxies And Through The Universe

Today, we know that most of the metals (70-90%) and nearly half of the baryons are unaccounted for, predicted to be in a hot phase (0.5-10E6 K) that is accessible to X-ray observations. X-ray absorption lines, mainly O VII, are detected in sight lines through the Milky Way halo, helping to define the density and mass of the gas at the virial temperature of about 2E6 K. No extragalactic sight lines show absorption, consistent with theoretical predictions. After Athena, there will be more high-quality sight lines through the Milky Way hot halo and the first extragalactic absorption systems should have been detected, through galaxy halos at impact parameters less than 100 kpc. Lynx will usher in a new age in that it will be able to detect absorption from galaxy halos out to the virial radius, and in the surrounding large-scale structures, the Cosmic Web absorption should be detectable to z&sim1, if not beyond. It will also measure the dynamics of the Milky Way hot halo &mdash rotation, turbulence, and accretion/outflow, thereby providing powerful constraints on galaxy formation and evolution.

Probing the Lifecycle of Nearby Galactic Nuclei via High-resolution X-ray Imaging and Spectroscopy

Galactic nuclei are believed to play a central role in galaxy formation and evolution. The lifecyle of their activity, however, remains greatly uncertain. I'll demonstrate how high-resolution X-ray imaging and spectroscopy of hot gas in galactic bulge regions can be used to probe AGN activity in the recent past (within one million year or so). We have developed an AGN relic model, which accounts for both photo-ionization and non-equilibrium evolution of distributed hot gas and allows for comparison with spatially resolved X-ray spectroscopic data. As an example, I'll describe the application of this model to the analysis of the Chandra and XMM-Newton X-ray data of the M31 central region, which is quiescent in both AGN and star formation, but shows strong indications for recent AGN activity. The diffuse X-ray emission from the high spatial resolution Chandra observations reveals interesting substructures with radio counterparts, as well as an overall negative temperature gradient in the diffuse hot gas distrbution. The X-ray grating spectra from the XMM-Newton observations show enhanced forbidden lines of He-like Oxygen, Neon, and Nitrogen Kalpha triplets, as well as signatures for multi-temperature hot gas. We find that these results can be well interpreted by the AGN-relic model, suggesting that galaxy is a bright AGN about 0.4 Myrs ago. In addition, we have also found evidence for resonance scattering effects, which broaden the spatial distribution of the relevant line emission and provide a sensitive probe of the hot gas turbulent motion. This application demonstrates the power of the spatially-resolved X-ray spectroscopy, as will be provided more effectively by Lynx, in our understanding of the recurrence history or frequency of AGN and galaxy feedback in general.

The XMM-Newton View of the z<0.5 Warm Hot Intergalactic Medium

We present preliminary results from the whole 1.6 Ms XMM-Newton observation of the z&sim0.5 Blazar 1ES 1553+113. The final 1.6 Ms spectrum of 1ES 1553+113 has reached a 90% sensitivity to 4 mA absorption line equivalent width. In the XMM-Newton and Chandra grating archives such sensitivities are reached only in the spectra of the brightest blazar in the Universe, Mkn421, which however explores a line-of-sight pathlength >10 times shorter than that seen against 1ES 1553+113. According to theoretical predictions between 3.1-6.2 WHIM OVII Ka absorbers should have been detected down to these sensitivities and up to such pathlengths along the line of sight to 1ES 1553+113. However, the RGS spectrum of 1ES 1553+113, which clearly detects several of the expected Galactic absorption lines down to such sensitivities and hints to a bunch of even weaker Galactic transitions, shows only two >3-sigma absorption lines potentially identifiable with intervening OVII Ka WHIM (and two more identifiable with NVII at the same redshift of one of the two OVII Ka, and NeX Ka at a different redshift). This invalidates the most optimistic predictions at >4-sigma confidence level and questions the most conservative ones, opening a number of questions that desperately need to be properly investigated and possibly addressed, both theoretically and observationally, before the advent of the next generation of high-resolution X-ray spectrometers, like Athena and/or Linx.

High Resolution Imaging and Gratings in Concert: An abundance of cutting edge science that can only be done with Lynx

Compared to the currently funded X-ray missions, Lynx offers two distinct advantages: high resolution imaging and high resolution gratings spectroscopy. High resolution imaging is necessary for studying stellar populations and diffuse plasma structures within crowded star forming regions, supernova remnants, and the Galactic Center. However, gratings spectroscopy is an often overlooked advantage of the Lynx mission. X-ray observations have a unique capability for studying metals in all phases, whether in cold neutral gas or hot ionized plasmas. In particular, carbon and oxygen &mdash the two most abundant metals &mdash have spectral features in the soft X-ray, where micro-calorimeters lose both sensitivity and resolving power. I will review the primary instrumentation limitations, avenues for solving them with Lynx, and the science outcomes to be expected, from supermassive black holes to the interstellar medium.

Measuring the Dust and Gas Phases in the Milky Way

Since its emergence, high resolution X-ray spectroscopy has greatly impacted studies of properties of the gas phases in the interstellar medium (ISM). Resolving the O K, Fe L, and Ne K edge structures revealed how X-ray spectra are affected by absorption and exposed the physics of the cold, warm, ionized and hot phases of the ISM. Some studies identified signatures related to dust and molecular components but their true existence remained unclear. Studies of higher Z edges such as Mg K, Si K and S K in contrast indicate dominant dust signatures in the edge structure. In a recent survey of the Si K edge in X-ray binaries located in the Galactic Bulge of the Milky Way we describe the edge using several components which include multiple edge functions, near edge absorption excesses from silicates in dust form, contributions from X-ray scattering optical depths as well a the presence of a variable warm absorber from atomic silicon. Some of the details require spectral resolutions of better than 3 eV, which even with the high energy gratings (HEG) onboard Chandra is already challenging. LYNX will provide us with ultra-resolved and high signal K edge structures from Mg, Si, S, Ar, Ca, Fe, and Ni allowing for the first time to effciently survey the Galactic plane and bulge in terms of identifying various dust species and their ratio to the gas phase. We present current results with Chandra on this topic and discuss them in the context of ISM physics and how this will be impacted by the power of LYNX.

Diverse Origin of Galactic Center X-ray Filaments

One of the most striking phenomena in the Galactic center are the numerous radio filamentary structures extending up to tens of parsecs. The radio polarization detection from the filaments point to a synchrotron origin with constant feeding of GeV electrons. In the X-ray band, about 20 elongated filaments have been discovered so far. They used to be interpreted as pulsar wind nebulae (PWNe). However, recent X-ray and radio observations have revealed a diverse origin of the X-ray filaments. In this talk, I am going to discuss the source nature of the three brightest X-ray filaments, as candidates for a PWN, a supernova remnant and cloud interaction, or a magnetic flux tube. Future deeper X-ray observation will hold the key to discover fainter X-ray filaments and to reveal their source nature.

An Exploration of The X-ray Emission from Normal Galaxies in the Early Universe with Lynx

Just above the detection limits of the unprecedented 7 Ms Chandra Deep Field-South, X-ray detected normal galaxies, powered primarily by X-ray binary (XRB) populations, now outnumber active galactic nuclei (AGN). As such, when we move into an era of even deeper X-ray surveys with Lynx, populations of distant (z>1.5) normal galaxies will be detected in large abundance, ready for exciting new studies of these populations and their cosmic evolution. In this talk, I will discuss recent results that show scaling relations between X-ray emission from low-mass XRBs (LMXBs) with stellar mass (LX/M) and high-mass XRBs (HMXBs) with star-formation rate (LX/SFR) evolve significantly out to z&sim2.5, such that LX(LMXB)/M&sim(1+z) 2-3 and LX(HMXB)/SFR&sim(1+z). I will put these findings into the context of theoretical population synthesis models, which indicate that the evolution XRBs can be attributed to global changes in the stellar ages and metallicities of galaxies in the Universe. Extending these models to higher redshifts suggest that the XRB emissivity of the Universe is expected to overpower AGN at z>5-8 and could play an important role in heating the intergalactic medium of the early Universe. I will conclude by discussing how future Lynx and multiwavelength observatories will extend constraints on XRB population evolution and how ground-based 21-cm experiment results could be paired with high-z X-ray galaxy studies to learn about these populations in the z&sim10-15 Universe.

Beyond Chandra: Accreting binary populations in the era of Lynx

X-ray binaries are our main tool for studying compact object populations and addressing key areas of Astrophysics, including the mass spectrum of compacts objects, the progenitors of gravitational waves and short gamma-ray bursts, and the preheating of the intergalactic medium in the epoch of reionization. Chandra and XMM-Newton have revolutionized our understanding of the X-ray binary populations and their dependence on the current and past star-forming activity of their host galaxy. They have even allowed us to set limits on parameters related to the formation and evolution of X-ray binaries. The unprecedented capabilities of Lynx and its synergy with the excellent suite of multi-wavelength observatories that are planned for the 2030s will enable the next quantum leap in studies of the X-ray binary populations in the nearby as well as the more distant Universe. We will discuss the prospects for studies of X-ray binary populations using a suite of end-to-end simulations of nearby as well as distant galaxies for different assumptions in X-ray binary population synthesis models. This way we will explore how Lynx will be able to constrain key parameters that determine the formation and evolution of X-ray binaries in the local Universe, by means of (a) resolved population studies in nearby galaxies, and (b) studies of ULXs and integrated X-ray emission of galaxies in wide area and deep surveys. Such constraints will have important implications for estimating the contribution of the host galaxy in studies of the gaseous component and the supermassive black-holes in high-z galaxies. In addition we will discuss how these results depend on key mission design parameters such as energy range and spatial resolution.

Lynx Data Analysis Challenges

Lynx promises to make remarkable progress in high-energy astrophysics instrumentation, with attendant improvements in data quality. As we learned during the development of Chandra, improvements in data quality require considerable thought on how to analyze the data. For instance, Chandra data has forced the migration of data analysis away from chi-square minimization, to the explicit use of Poisson likelihoods, and the common use of MCMC. It is illustrative to realize that wavdetect was developed just to apply to Chandra images. These techniques were developed and used mainly to obtain better and more robust astrophysical results. We anticipate a similar revolutionary advance with Lynx. We describe several examples of how Lynx data analysis will benefit from advanced techniques that are being currently developed or are anticipated. We consider issues like image deconvolution/segmentation/reconstruction, non-parametric flux estimation from calorimeter spectra, robust photometry in crowded fields, combining datasets, etc. We invite people to join us in an informal Working Group whose goal is to anticipate and game out upcoming data and analysis issues.

Observing Supernova Remnants with Lynx

X-ray observations of supernova remnants (SNRs) are a crucial means to investigate the mechanisms of SN explosions and the acceleration of particles to TeV energies by SN shocks. Since Chandra's first-light image of Cassiopeia A, modern X-ray facilities have revealed that SNRs have complex morphological and spectral characteristics that give important insights to the nature of explosions, progenitor stars, and their effects on their surroundings. In this presentation, I will review some of the major advances in SNR science in the past 18 years, focusing principally on results and discoveries from Chandra data. I will also discuss the open questions in the field and Lynx's role in addressing these outstanding issues.

Probing Mass Loss in Supernova Progenitors with Lynx

There is extensive evidence that the progenitors of some core collapse supernovae undergo extensive mass loss in the years leading up to core collapse. The imprint of the mass loss episodes may be seen in the spectra and dynamics of the expanding remnant, on timescales of years to centuries after the supernovae. In this talk, I will discuss how Lynx observations of young supernovae in the local Universe will allow us to understand mass loss mechanisms in massive stars. I will discuss how high resolution spectra will allow us to reconstruct the progenitor mass loss history and infer properties of massive stars in the late stages of their evolution.

Constraining the Dense Matter Equation of State with Lynx Observations of Globular Cluster X-ray Binaries

Extensive Chandra surveys have revealed that globular clusters are veritable treasure troves of X-ray faint compact binaries. Among them are the class of X-ray binaries with thermally-emitting "quiescent" neutron stars. Sensitive spectroscopic observations of such systems can provide limits on the mass-radius relation of neutron stars, thereby producing constraints on the pressure-density relation of matter at extreme densities. Effective studies of globular cluster sources require sub-arcsecond resolution, which among the future generation of X-ray missions only Lynx would provide. I will describe the prospects of dense matter equation of state measurements with Lynx as well as the wealth of secondary science that would be enabled by the same observations.

Massive molecular gas flows and AGN feedback in galaxy clusters

Powerful radio jets launched by a central supermassive black hole pump a substantial amount of energy into their host galaxies and cluster environment. This feedback from the central AGN is thought to suppress gas cooling and star formation at late times to regulate the growth of massive galaxies and their hot atmospheres. The most massive galaxies, located at the centres of rich clusters, often host substantial reserves of molecular gas, which are likely fuelling star formation and the black hole activity. I will review ALMA Early Science observations that resolve the structure of these molecular reservoirs and reveal massive filaments extending 5-15 kpc, which likely formed from gas cooling out of the surrounding hot atmosphere. I will present recent ALMA observations of extended molecular filaments in the Phoenix, PKS0745-191 and Abell 1795 central cluster galaxies, which are drawn up around giant radio bubbles. Smooth velocity gradients along these filaments are consistent with ordered molecular gas flows around the bubbles and show that radio jets interact with cold, dense molecular gas as well as the hot, diffuse intracluster medium. I will also discuss the exciting possibilities of combining ALMA and Lynx observations.

Stellar Coronal Studies at High Spectral Resolution with Lynx

The search for life around stars other than the Sun requires detailed knowledge of the host stars themselves. The star interacts with an exoplanet's environment through X-ray and extreme ultraviolet radiation and high energy particle fluxes, manifestations of the star's magnetic activity. Stellar winds can exhibit high particle enhancements during coronal mass ejections associated with energetic stellar flares. Deleterious effects on habitability by the star's magnetic activity are easy to imagine at the same time, some level of activity may also be necessary to catalyze life.

In this talk, I will review the status of some of the big questions in coronal physics: the coronal heating problem, unsolved since its identification in 1941 stellar dynamo theory, recently highlighted by Chandra studies of fully convective M dwarfs and the relationship between stellar winds and the confinement of plasma in coronal loops.

X-ray grating spectra obtained with Chandra and XMM-Newton reveal high electron densities and sharply peaked emission measure distributions in the coronae of active stars, unlike features observed in the solar corona. These spectra typically have long exposure times, and thus are limited to nearby, highly active stars, with poor sampling of accreting systems, lower mass dwarfs and older, slow rotators like the Sun. Even at the highest spectral resolution of the Chandra gratings, many diagnostic line ratios are blended, velocity measurements are rare, and equilibrium tests are not possible. Lynx offers large improvements in spectral resolution and throughput. I will discuss new types of measurements that Lynx will enable, including velocities and tests of thermal equilibration, and how these will inform our understanding of stellar coronae.

What can Lynx do for coronal physics?

The solar corona shows us a range of magnetic and plasma phenomena that are also going to be present on other stars. Stars allow us to study what happens to these phenomena when driven to extremes of magnetic activity, or when displayed by stars of different underlying fundamental parameters. We look at these phenomena, from wave process to reconnection, from thermal instabilities and oscillations to flares and CMEs, and highlight the physical processes that a next generation X-ray mission such as Lynx can hope to observe and help understand.

Observational signatures of self-consistent 3D CME-flare models

We are examining in a numerical study the relation between CMEs and flares in different stellar activity regimes. The computational tools used to achieve this task are BATS-R-US on the magnetohydrodynamic end and (implicit) iPIC on the kinetic end. A bidirectional coupling between the two codes has been implemented in order to examine the heating, acceleration and emission mechanisms linked to stellar CMEs and flares. In particular, magnetic reconnection, that is a kinetic process, takes place at the coronal loop apex. As a result of the reconnection, electrons get accelerated and as they descend with high velocities from the sparse corona they collide with the denser chromospheric material and emit in hard X-rays creating two flaring sites at the footpoints of the loop. Three distinct X-ray bright regions are the observational signatures of the whole process, namely, the loop apex (soft X-rays) and its footpoints (hard X-rays). We examine how numerical modelling of these processes coupled with observations of stars with Lynx might be used to provide new physical insights into stellar flares and CMEs.

Observations of Stellar Clusters From Chandra to Lynx.

In this presentation I discuss what we have learned from Chandra observations of nearby young stellar clusters such as NGC 1333, IC 348, The Orion Nebular Cluster and more distant Clusters such as Eta Carina and NGC3603. Then I examine what the additional capabilities of Lynx would imply for similar observations. The net result is that Lynx should be able to make complete samples of Young stellar clusters out to 2 kpc. This is similar to the grasp of Gaia and JWST. I close with a discussion of problems which will exist in the 2030's which Lynx may be able to address.

High Resolution X-ray Spectroscopy and Star Formation: HETG Observations of the Pre-Main Sequence Stellar Cluster IC 348

We present Chandra High Energy Transmission Grating (HETG) observations of the &sim3 Myr old pre-main sequence (pre-MS) stellar cluster IC 348. With 300-400 cluster members at a distance of &sim300 pc, IC 348 is an ideal target to observe a large number of X-ray sources in a single pointing and is thus an extremely efficient use of Chandra-HETG. High resolution X-ray spectroscopy offers a means to investigate detailed spectral characteristic of X-ray emitting plasmas and their surrounding environments. We present our initial findings where we compare X-ray spectral signatures (e.g., luminosity, temperature, column density, abundance) of the X-ray brightest pre-MS stars in IC 348 with spectral type, multiwavelength signatures of accretion, and the presence of circumstellar disks at multiple stages of pre-MS stellar evolution. Assuming all IC 348 members formed from the same primordial molecular cloud, any disparity between coronal abundances of individual members, as constrained by the identification and strength of emission lines, will constrain the source(s) of coronal chemical evolution at a stage of pre-MS evolution vital to the formation of planets. Chandra-HETG observations of pre-MS stellar clusters like those presented here for IC 348 are essential motivation for the inclusion of gratings on Lynx, the next generation X-ray telescope. For HETG observations of nearby pre-MS stellar clusters, Chandra's sensitivity is capable of producing high signal-to-noise ratio spectra for only the X-ray brightest cluster members in a reasonable exposure time. With the significant increase in sensitivity of Lynx, observations of crowded star-forming regions can produce more high resolution spectra at a fraction of the exposure time resulting in an unprecedented view of pre-MS stellar evolution.

The roles of host star XUV emission on exoplanet habitability and physical processes in stellar coronae and the ISM

With 10,000 resolution, 0.5 arcsec angular resolution, and high throughput, the proposed Lynx XUV observatory could provide valuable insights into major issues in astrophysics. In this talk, I will describe the many ways in which observations with Lynx are needed to understand the evolution of exoplanet atmospheres. I will also mention some important discoveries that Lynx could make concerning coronal physics and the interstellar medium.

The development of organic chemistry in protoplanetary disks likely begins with the host star's XUV emission photoionizing hydrogen to produce H2+ and H3+. The photoionization is primarily at EUV wavelengths which are mostly unobservable through the interstellar medium. With its high spectral resolution, Lynx could separate X-ray emission lines from continua leading to more accurate emission measure distributions from which the unobservable EUV emission can be calculated. As protoplanets evolve, the X-ray and EUV emission of host stars photoionize and heat their outer atmospheres, leading to hydrodynamic mass loss, which may remove the entire atmosphere. EUV and UV photons also photodissociate important molecules in exoplanet atmospheres. With a resolution of 30 km/s, Lynx could measure the accretion flow of protoplanetary disk gas onto the host star and search for coronal mass ejections (CMEs) that can destroy ozone in planetary atmospheres, thereby sterilizing the planetary surface. CMEs may be detected by their transient X-ray absorption during flares.

Lynx should search for Doppler shifts in coronal emission lines to study upflows and downflows during stellar flares and outflowing hot gas emission from protostellar disks and perhaps even quiescent stellar coronae. Lynx spectra will be useful for measuring coronal electron densities from the He-like triplet lines, including the blended Ne IX lines, and lines of Fe XXI and Fe XXII. Lynx spectra may be able to study departures from collisional ionization equilibrium at quiescent times and during flares. Measuring fluxes of weak emission lines will be needed to study a wide range of ionization stages. Spectra of transition region lines such as He II 30.4 nm can test for theimportance of ambipolar diffusion. Finally, high-resolution spectra of the He II 30.4 nm resonance line and the ionization edge of He II at 22.8 nm could measure the abundance of ionized helium in the local interstellar medium.

From ChanPlaNS to Lynx

In stellar evolution one of the shortest evolutionary phases is the planetary nebula (PN) phase. The phase exists for tens of thousands of years after the mass lost during the asymptotic giant branch phase is sculpted into a nebular shell and the emerging hot stellar core, which ionizes the shell, fades towards the white dwarf cooling track. The Chandra Planetary Nebulae Survey (ChanPlaNS) identified an even shorter phase that pertains to the period when the nebula is filled with a X-ray emitting plasma called a hot bubble. This hot bubble phase exists early in the lifetime of the PN and is the result of the initial violent collision of stellar winds that shapes the nebula. A surprising result made by Chandra are the homogeneous and unexpectedly cool plasma temperatures of these hot bubbles. It is believed that heat conduction across the nebular-hot bubble interface leads to mixing of these two plasmas and cooling of the hot bubble to the observed range of 1-3 MK. The heat conduction process is uncertain because these hot bubbles are chemically enriched and difficult to detect. Their cool temperatures (1-3 MK) require sensitivity to soft X-ray photons (E<1 keV) and only one PN has a high enough flux to obtain the high quality grating spectrum required to study its chemistry. Future gratings observations by Athena+ and/or Hitomi 2 will help establish the chemical distribution amongst hot bubbles in PNe, but high sensitivity and high spatial resolution imaging, like that proposed for Lynx, is essential in our study of these hot bubbles. Lynx will allow us to map the spatial distribution of the hot bubble plasma properties. This spatial information provides a key test of the heat conduction process. Ultimately, Lynx will allow us to realize the full potential of these wonderful astrophysical plasma laboratories.

Winds of Massive Stars: Line Profiles and Variability

The X-ray line profiles from the winds of massive stars are determined by the wind structure and dynamics. Stellar winds, however, are not necessarily uniform or constant, as is seen in features migrating through UV and optical lines, interpreted as co-rotating interaction regions (CIR). We have searched for line variability in the O-star, zeta Puppis (among others), with suggestive but unconvincing results. With greater sensitivity and high spectral resolution, we would be able to determine if there are X-ray counterparts to CIR and refine our understanding of plasma heating in winds and the use of X-ray emission lines as diagnostics of mass loss.

Chandra Implications for Future Calorimeter Observations of Starburst Feedback

Lynx promises to make critical, and potentially game-changing, improvements to our understanding of supernova-driven feedback in starburst galaxies by resolving down to the arcsecond scales of structure seen in deep Chandra observations and at other wavelengths. A first step at assessing this would be to produce diffuse X-ray flux maps for nearby galaxies to derive the temperature and abundance distribution based on Chandra observations. Here we describe our analysis of the Chandra data for nearby starburst galaxies and the use of these data to simulate future Lynx and Athena observations. Our main goals will be to determine how well Lynx will be able to constrain stellar feedback in nearby galaxies and to assess design requirements for Lynx to constrain the energetics of stellar feedback, map the dispersion of metals into the ISM and IGM and to resolve the interaction of the superwind fluid with swept-up ISM on the scales of filaments and shocks. Additionally Lynx holds the promise of resolving extranuclear diffuse emission at distances &sim10× further than Athena, and have the sensitivity to detect the hot superwind fluid directly in some systems, e.g., by detecting H-like Fe-K emission.

Spatially Resolving AGN in Transitioning Galaxies with Chandra

The role of AGN in the transition of galaxies from actively star forming to quiescence is still not fully understood, and is particularly challenging to study for AGN that are weak or hidden by obscuration or host dilution. The high spatial resolution of Chandra is critical in separating the AGN from the host emission in these complex systems. The Shocked Post-starburst Galaxy Survey (SPOGs) selects quenching galaxies still containing shocks and molecular gas based on their optical line diagnostics. SPOGs catches galaxies at an earlier stage of transition than classical post-starburst selection criteria, providing a crucial window into the role of AGN in this early phase of transformation. We recently obtained a deep Chandra observation of the proto-typical SPOG NGC1266 combined with a NuSTAR observation as well as Chandra observations of a set of SPOGs showing enhanced IR emission relative to their molecular content. I will discuss the insights obtained about the characteristics of AGN in SPOGs and the implications on the role of AGN in galaxy transitions.

Athena: Science Prospects and Mission Status

The Athena X-ray observatory has been conceived to study the Hot and Energetic Universe. The principal areas of study will include studies of the physical and chemical evolution of the hot baryonic component in the Universe, the physics of galaxy clusters, phenomena generated around black holes and feedback on all mass scales. In this talk, I will detail the science objectives of Athena and provide an update on the mission's status in its present study phase, with the goal of being proposed for 'adoption' around 2019/20.

AGN Feedback Lessons from Chandra Observations of Clusters of Galaxies

I will review a few of the key insights about how AGN feedback works to regulate itself in massive galaxies, yielded primarily through observations with the Chandra X-ray Observatory. I will present a framework that can be applied to lower-mass systems. I will show and provide a recipe for predicting the halo properties of the most luminous X-ray systems at range of galaxy mass scales.

Analyzing Cocoon Shocks in Radio Galaxies: Predictions for Lynx

Chandra observations have highlighted the role of radio outbursts from cluster central galaxies in limiting star formation in cool cores. In addition to being hosted by the central galaxy of a cool core cluster, the FRII radio galaxy Cygnus A is the archetype of powerful radio galaxies. I will present results from a detailed study of the cocoon shocks of Cygnus A using 2 Msec of Chandra observations. The AGN outburst has an age close to 20 Myr and a mean power of about 10 46 erg/sec. Weak shock strengths are detected over the majority of the cocoon, implying a uniform pressure in the bulk of the radio lobes. Estimated hotspot pressures agree with synchrotron self-Compton models, while determined jet and lobe properties are broadly consistent with FRII models. These results required close to the deepest feasible Chandra observations, but similar results could be achieved routinely with the high throughput of Lynx. Furthermore, the high resolution X-ray spectroscopy of Lynx will enable us to examine issues currently inaccessible with Chandra, as I will demonstrate with simulations.

X-ray insights into the earliest stages of a radio source evolution

Compact Symmetric Objects (CSOs) are thought to be among the progenitors of large-scale radio galaxies. They show radio features typically observed in large-scale radio galaxies (jets, lobes, hot spots), but contained within the central 1 kpc region of the host galaxy. Because the CSOs are symmetric and not affected by beaming, their linear radio size can be translated into the source age if one measures the expansion velocity of the radio source. However, if the jet expansion is disturbed, e.g. by a dense interstellar medium (ISM), the ages derived this way may be biased. Until now we did not have means to discriminate between confined and non-confined radio sources. We present our X-ray studies of CSOs performed with Chandra and XMM-Newton. For the first time, the data reveal the evidence in favor of the idea that in a sub-population of CSOs the radio jets may be confined by a dense ISM. Thus, the CSO kinematic ages may be underestimated. We discuss the implications of our results on the high energy emission models of CSOs, the earliest stages of a radio source evolution, jet-galaxy co-evolution, and advances that are expected in the CSO field thanks to Lynx.

Observing the First Galaxy Clusters with Lynx

In recent years, the combination of SZ surveys and X-ray follow-up has led to rapid advances in our understanding of galaxy clusters and their evolution over the past

10 billion years. I will summarize the progress we have made in a small amount of time, exploiting the synergy of the South Pole Telescope and Chandra X-ray Observatory. In the next decade, the number of known clusters at high redshift is poised to grow by several orders of magnitude, with X-ray surveys by eRosita, SZ surveys by SPT-3G, Advanced ACTPol, and CMB-S4, and OIR surveys by LSST, Euclid, and WFIRST. The most distant of these systems will be extremely faint and soft &mdash we currently have no ability to follow these systems up in the X-ray. I will summarize several exciting programs that will be possible with Lynx in combination with these next-generation surveys, in particular highlighting the need for soft sensitivity and high angular resolution in any next generation X-ray mission.

Cluster AGN Topography Survey

The nature of SMBH-galaxy co-evolution remains one of the major outstanding problems in modern astronomy. A natural link exists between the fuelling of star formation in galaxies and SMBH growth, namely the availability of cold gas. The cluster environment influences gas reservoirs through processes such as ram-pressure stripping, evaporation, starvation, and tidal effects of the cluster potential. The relative importance of these processes depends on both the position within, the mass of, and the redshift of the host galaxy cluster. As such, comparisons of AGN populations between dense cluster environments and the field provides valuable additional constraints on fuelling processes, isolating the connection between SMBH and galaxy growth.

Our goal is to investigate how these environmental processes influence the evolution of the cluster AGN population. To this end, we have launched a massive, multi-wavelength campaign anchored by Chandra data. Chandra's excellent spatial resolution is allowing us to explore the X-ray AGN population in more than 500 massive galaxy clusters from z=0 to z&sim1.8. An X-ray AGN survey is crucial for such a study as it suffers little contamination, has low absorption bias, and provides direct access to most of the accretion power in the Universe. We expect to identify &sim40,000 X-ray AGN and will trace the evolution of the cluster AGN population across cosmic time and as a function of cluster mass. Using VLA FIRST data, we are performing a parallel analysis of the radio AGN population in a subset of these clusters, providing a complementary probe into the evolution of kinetic mode AGN activity. Furthermore, we have obtained deep VIMOS optical spectroscopy of 7 of these clusters, allowing us to extend our analysis to a background free investigation of the cluster AGN population as a function of stellar mass and star formation rates.

Looking to the future, Lynx's high spatial resolution, wide field of view, and large effective area at soft energies, uniquely enable an extension of this study out to both the highest redshift and lowest mass clusters. For the first time we will be able to investigate the complete story of AGN evolution in dense environments, from early-times, through the peaks of AGN and SF activity, to the present day.

The Impact of Environment on Galaxy Cluster Outskirts

A significant fraction of the mass of galaxy clusters is located in the outskirts of clusters, outside of their relatively well-studied central regions. Therefore, understanding the structure and ICM physics of clusters in this region, and how they depend on cluster dynamical state and local environment, is important for doing precision cosmology with clusters. Despite this fact, cluster outskirts remain relatively poorly understood. For example, entropy profiles of the ICM generally lie below what would be expected from simple gravitational collapse models of structure formation, possibly due to the presence of small unresolved gas clumps, although the clumping factors predicted by simulations are generally too small to explain observations. I will discuss what we know about the virialization region of clusters, with a focus on our recent work on merging clusters and the connection between cluster outskirts and large scale structure filaments. I will also discuss how upcoming and proposed X-ray missions such as Lynx can contribute to this important area of study.

Lessons learned from Chandra AGN surveys and what's next

The extragalactic surveys performed by Chandra over the past 18 years have shed light on AGN demographics and greatly shaped our current understanding of black hole and galaxy co-evolution. The combination of deep and large area X-ray studies, which probe a wide range of the luminosity-redshift space, and the associated multiwavelength data are the keys to the achieved progress.

In this talk, I will present the work performed primarily with the Chandra COSMOS Legacy survey as well as deeper and wider surveys, using both detected sources and stacking analysis of non-detections. I will focus on the relation between BH accretion and star formation, and SMBH growth in the high redshift Universe, including what we have learned about black hole seeds from nearby intermediate mass black holes. I will also present how archival X-ray and multiwavelength data allows us to derive cosmological parameters using AGN/QSO as standard candels. Finally, I will present what we will be able to learn on these subjects with Lynx and complementary multiwavelength data in the 2030s.

Harder, Better, Faster, Stronger: Mapping Extended Line Emission from AGN above 1 keV

Active galactic nuclei (AGN) tend to ionize gas in their host galaxies via photoionization and kinematic outflows, creating extended X-ray signatures of feedback that Chandra can spatially resolve. Although the bulk of this emission occurs below 1 keV, in recent years deep Chandra observations of nearby AGN have shown extended harder emission both in lines and continuum, which has fundamentally challenged our conception of the nature of hard X-rays from AGN. We will discuss ways in which Chandra's ability to spatially resolve harder X-ray emission presents opportunities for deep observations in anticipation of Lynx's unprecedented sensitivity and energy resolution, in an era where the ACIS contaminant buildup makes this regime a natural focus.

Prospects for Large-scale Jets with Lynx

High-resolution instruments like the VLA, HST, and Chandra have allowed us to map the structure and SED evolution of extragalactic jets on the kpc scale where they interact with the galactic and intergalactic environment, yet the physical description of these jets remains mysterious. Among things remaining to be settled are the particle makeup of these jets, their velocity profiles and total energy content, and in many cases even the radiation mechanism. One of the great discoveries by Chandra has been the 'anomalous' class of X-ray bright jets, for which the X-ray mechanism is still unsettled. I will describe some recent developments in the study of resolved X-ray detected AGN jets, and discuss how a successor to Chandra will be critical to finally solving the X-ray origin problem in large-scale jets.

Physical properties of the highest redshift supermassive black holes

I will discuss the current understanding of key physical properties of some of the first generation of growing supermassive black holes (SMBHs). This includes their accretion rates and history, their host galaxies, and the large-scale environments that enable their emergence about a billion years after the big-bang. The available multi-wavelength data is consistent with Eddington-limited, radiatively efficient accretion that had to proceed almost continuously since very early epochs. New ALMA data confirms high SFRs and gas content in the host galaxies, and moreover a high fraction of companion, interacting galaxies, separated by

10-50 kpc. This clearly support the idea that the first generation of luminous SMBHs grew in overdense environments, and that major mergers are important drivers for rapid early SMBH and host galaxy growth, although other fueling processes and accretion modes may still be required. Current X-ray surveys cannot access the lower-mass counterparts of these rare massive quasars, which would elucidate the earliest stages of BH formation and growth. Such low-mass nuclear BHs will be the prime targets of the deepest surveys foreseen for Lynx, within the theme of "First Accretion Light".

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Recent observations of sulfur containing species (SO2, SO, OCS, and H2SO4) in Venus’ mesosphere have generated controversy and great interest in the scientific community. These observations revealed unexpected spatial patterns and spatial/temporal variability that have not been satisfactorily explained by models. Sulfur oxide chemistry on Venus is closely linked to the global-scale cloud and haze layers, which are composed primarily of concentrated sulfuric acid. Sulfur oxide observations provide therefore important insight into the on-going chemical evolution of Venus’ atmosphere, atmospheric dynamics, and possible volcanism.

This paper is the first of a series of two investigating the SO2 and SO variability in the Venus atmosphere. This first part of the study will focus on the vertical distribution of SO2, considering mostly observations performed by instruments and techniques providing accurate vertical information. This comprises instruments in space (SPICAV/SOIR suite on board Venus Express) and Earth-based instruments (JCMT). The most noticeable feature of the vertical profile of the SO2 abundance in the Venus atmosphere is the presence of an inversion layer located at about 70–75 km, with VMRs increasing above. The observations presented in this compilation indicate that at least one other significant sulfur reservoir (in addition to SO2 and SO) must be present throughout the 70–100 km altitude region to explain the inversion in the SO2 vertical profile. No photochemical model has an explanation for this behaviour. GCM modelling indicates that dynamics may play an important role in generating an inflection point at 75 km altitude but does not provide a definitive explanation of the source of the inflection at all local times or latitudes

The current study has been carried out within the frame of the International Space Science Institute (ISSI) International Team entitled ‘SO2 variability in the Venus atmosphere’.