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This also means that the resonant peaks shown in Figure 5 are not the highest that are possible for P,2—=0.03. but the resonance is still quite pronounced for the parameter values chosen for Figure 5 | This also means that the resonant peaks shown in Figure 5 are not the highest that are possible for $P_{\alpha_L}=0.03$, but the resonance is still quite pronounced for the parameter values chosen for Figure 5. |
Our principal conclusion is Chat. at least in this simple dyvniamo model. resonance will always be found. provided the right parameter values are chosen. | Our principal conclusion is that, at least in this simple dynamo model, resonance will always be found, provided the right parameter values are chosen. |
Furthermore. it occurs for parameter choices that are plausible for the sun. | Furthermore, it occurs for parameter choices that are plausible for the sun. |
It occurs in the lowest laver of the model. where the meridional How toward the equator is most likely to match. or nearly match. the propagation speed of the forcing imposed at the top. corresponding to the photosphere on ihe sun. | It occurs in the lowest layer of the model, where the meridional flow toward the equator is most likely to match, or nearly match, the propagation speed of the forcing imposed at the top, corresponding to the photosphere on the sun. |
And the effect is lareea factor of 10-100 zumplilication of toroidal field compared to the case where the parameter values are far from the ones for which resonance is predicted. | And the effect is large–a factor of 10-100 amplification of toroidal field compared to the case where the parameter values are far from the ones for which resonance is predicted. |
We acknowledge that this large a difference in amplitude for different parameter values is partly due to the model being kinematic. | We acknowledge that this large a difference in amplitude for different parameter values is partly due to the model being kinematic. |
IE XD. [orces were included. the peaks would almost surely be smaller. | If $jXB$ forces were included, the peaks would almost surely be smaller. |
The solar convection zone is a spherical shell. not an infinite laver. so how well should | The solar convection zone is a spherical shell, not an infinite layer, so how well should |
Although the mismatch between the a...) profiles andthe Geos profiles in the higher-mass simulations (see Figure +)) suggests the presence of non-local transport. it does not tell us whether these simulations reach quasi-steady states or not. | Although the mismatch between the $\alpha_{\rm total}$ profiles andthe $\alpha_{\rm cool}$ profiles in the higher-mass simulations (see Figure \ref{fig:Ms1alpha}) ) suggests the presence of non-local transport, it does not tell us whether these simulations reach quasi-steady states or not. |
To identify how quasi-steady the dises are. the discs’ temperature protiles and Toomre instability protiles are averaged over the tinal 13 ORPs. | To identify how quasi-steady the discs are, the discs' temperature profiles and Toomre instability profiles are averaged over the final 13 ORPs. |
The standard deviation about this mean is then measured. and the normalised quantities er/7 and σος) are calculated for each radius (Figure 51), | The standard deviation about this mean is then measured, and the normalised quantities $\sigma_{\rm T}/T$ and $\sigma_{\rm Q}/Q$ are calculated for each radius (Figure \ref{fig:TQ_Ms1}) ). |
This shows the deviation of the disc from quasi-steady. thermal equilibrium (through op /77) and its deviation from a marginally-stable. self-regulated state (through me /(Q). | This shows the deviation of the disc from quasi-steady, thermal equilibrium (through $\sigma_{\rm T}/T$ ) and its deviation from a marginally-stable, self-regulated state (through $\sigma_{\rm Q}/Q$ ). |
Simulation | (diui0.25. solid line in Figure 59) shows the lowest temperature deviation. maintaining thermal balance to within around except in the outer regions. where this is mainly due to the reduced value of 7. | Simulation 1 $q_{\rm init}=0.25$, solid line in Figure \ref{fig:TQ_Ms1}) ) shows the lowest temperature deviation, maintaining thermal balance to within around except in the outer regions, where this is mainly due to the reduced value of $T$. |
A deviation of 1 K from a mean of 10 K will be more significant than from a mean of 100 K. This is also true for diui;=0.5 (dotted line in Figure 59). although the amplitude increases further at larger radii. | A deviation of 1 K from a mean of 10 K will be more significant than from a mean of 100 K. This is also true for $q_{\rm init}=0.5$ (dotted line in Figure \ref{fig:TQ_Ms1}) ), although the amplitude increases further at larger radii. |
The lower-mass simulations (diii 0.5) are therefore not only local. but also settle into long-lived. quasi-steady« states. | The lower-mass simulations $q_{\rm init} < 0.5$ ) are therefore not only local, but also settle into long-lived, quasi-steady states. |
The temperature profiles for the high-mass (fini25 0.5) dises (dashed and dash-dot lines in Figure 59) show significant variation (varying by as much as in the worst case). illustrating that these discs not only have non-local transport. but also do not attain well-detined. long-lived quasi-steady states. | The temperature profiles for the high-mass $q_{\rm init} > 0.5$ ) discs (dashed and dash-dot lines in Figure \ref{fig:TQ_Ms1}) ) show significant variation (varying by as much as in the worst case), illustrating that these discs not only have non-local transport, but also do not attain well-defined, long-lived quasi-steady states. |
This implies that in these dises - at any given location - there will be periods when the dissipation rate exceeds the local cooling rate (causing the temperature to rise) followed by a period when the cooling rate dominates. | This implies that in these discs - at any given location - there will be periods when the dissipation rate exceeds the local cooling rate (causing the temperature to rise) followed by a period when the cooling rate dominates. |
This is presumably inherently linked to the global nature of the energy transport in these simulations. | This is presumably inherently linked to the global nature of the energy transport in these simulations. |
Energy is being transported non-locally. and is hence not being generated and dissipated at the same location. and therefore it is not possible for the local heating and cooling rates to balance at all locations in the disc. | Energy is being transported non-locally, and is hence not being generated and dissipated at the same location, and therefore it is not possible for the local heating and cooling rates to balance at all locations in the disc. |
Figure S also shows deviations from uniform €. with again the lower-mass dises showing the lowest deviation in the inner SO au. averaging around105c. | Figure \ref{fig:TQ_Ms1} also shows deviations from uniform $Q$, with again the lower-mass discs showing the lowest deviation in the inner 50 au, averaging around. |
. Simulations 3 4 (gii7 0.5) again vary much more significantly. peaking at around40%. | Simulations 3 4 $q_{\rm init} > 0.5$ ) again vary much more significantly, peaking at around. |
. These results show that for diui0.5 a dise is unable to settle into a long-lived. marginally-stable. self-gravitating state. | These results show that for $q_{\rm init} > 0.5$ a disc is unable to settle into a long-lived, marginally-stable, self-gravitating state. |
Although the above suggests that there is non-local transport in the higher-mass discs. we have not yet convincingly shown that thisis indeed the case. | Although the above suggests that there is non-local transport in the higher-mass discs, we have not yet convincingly shown that thisis indeed the case. |
One way to do this is to compare the pattern speed of the dominant spiral mode. £2,. with the angular speed of the dise material itself. OQ. | One way to do this is to compare the pattern speed of the dominant spiral mode, $\Omega_p$, with the angular speed of the disc material itself, $\Omega$. |
As shown by ?.. transport through gravitational instability can only be described in viscous terms when (23,= © | As shown by \citet{Balbus1999}, transport through gravitational instability can only be described in viscous terms when $\Omega_p = \Omega$ . |
When 2,=©. the non-local transport terms become more significant. | When $\Omega_p \ne \Omega$, the non-local transport terms become more significant. |
Waves producing non-local transport therefore have a pattern speed that deviates significantly from corotation (22).. | Waves producing non-local transport therefore have a pattern speed that deviates significantly from corotation \citep{Balbus1999,Cossins2008}. . |
Equivalently. the non-local transport fraction £ must deviate significantly from zero (2).. where [O,Q| can be calculated from the dispersion relation for finite thiekness Kepleriandises (2? | Equivalently, the non-local transport fraction $\xi$ must deviate significantly from zero \citep{Cossins2008}, , where $|\Omega_p - \Omega|$ can be calculated from the dispersion relation for finite thickness Kepleriandiscs \citep{Bertin2000,Cossins2008}
|
neutral Na is depleted at high altitudes. prestunably by ionization. below an arbitrary pressure. | neutral Na is depleted at high altitudes, presumably by ionization, below an arbitrary pressure. |
For the calculation. we ft various simplified T-P profiles in which the temperature is taken as a linear function of the altitude aud pressure is calculated using lyvdrostatic equilibrium. | For the calculation, we fit various simplified T-P profiles in which the temperature is taken as a linear function of the altitude and pressure is calculated using hydrostatic equilibrium. |
Note that it is not specifically the T-P profile shape itself which is important iu our fitting of the data. but rather the different teniperature and pressure regimes nuposed bv the observations. | Note that it is not specifically the T-P profile shape itself which is important in our fitting of the data, but rather the different temperature and pressure regimes imposed by the observations. |
At our signal-to-noise levels. the general variation of the T-P profile in between the different atinosplierie regimes we find (hot at high pressure. cool at middle pressures. aud hot at low pressures) is only loosely coustrained. | At our signal-to-noise levels, the general variation of the T-P profile in between the different atmospheric regimes we find (hot at high pressure, cool at middle pressures, and hot at low pressures) is only loosely constrained. |
A more complicated T-P profile would add undesired. acditional paralucters into the fit. leading us to choose as simple a T-P profile as possible. | A more complicated T-P profile would add undesired additional parameters into the fit, leading us to choose as simple a T-P profile as possible. |
We found it necessary to have a nüununn of three variable points in the T-P. profile. | We found it necessary to have a minimum of three variable points in the T-P profile. |
The temperature as a function of the altitude. Z(:). is thus composed of two linear fictions below and above the altitude z,,. where ty Is the altitude of the miuiinuua temperature. ων | The temperature as a function of the altitude, $T(z)$, is thus composed of two linear functions below and above the altitude $z_{m}$, where $z_{m}$ is the altitude of the minimum temperature, $T_{m}$. |
We took the lower reference altitude. τς, for the altitude corresponding to the mean planetary radius (Ixuutsouetal. 2007).. aud the altitude τε is the fixed. position. at lieh altitude. needed to defiue the second. part of the T(:) linear function. | We took the lower reference altitude, $z_{s}$, for the altitude corresponding to the mean planetary radius \citep{Knut}, and the altitude $z_{th}$ is the fixed position, at high altitude, needed to define the second part of the $T(z)$ linear function. |
We used τμ corresponding to the altitude of the absorption depth of which corresponds to the mean absorption depth iu the core of the sodimm doublet. | We used $z_{th}$ corresponding to the altitude of the absorption depth of, which corresponds to the mean absorption depth in the core of the sodium doublet. |
The temperature aud pressure at the altitude τε. the altitude z,, of nininmua temperature T,,,. aud Ty), ave the five free paraincters define the T- profile. | The temperature and pressure at the altitude $z_{s}$, the altitude $z_{m}$ of minimum temperature $T_{m}$, and $T_{th}$ are the five free parameters defining the T-P profile. |
The κοπα abundance in the lower aud the upper atinosphere are the two other free parameters iu the fit. | The sodium abundance in the lower and the upper atmosphere are the two other free parameters in the fit. |
Tn our coudensation fit we found: (1) a high temperature of 2.200260 IK at 3325 mbar (500 I) mainly coustrained by Bavleigh scatteriug(Lecavelieretal.2008b) seen at wavelengths shorter than ~5.000A. (2) a low temperature of 12025190 Is at 0.6340.27 mbar needed to explain the plateau at (51.500. kia) above the minium absorption level as a Na abundance change due to condensation. and finally (3) a high altitude (73.500 Iu) hieh temperature of IN. at 0.0063250.0023 mbar needed to fit the strong peak 1.770228aud narrow width of the cores of the Na I doublet. | In our condensation fit we found: (1) a high temperature of $\pm$ 260 K at $\pm$ 5 mbar $\sim$ 500 km) mainly constrained by Rayleigh scattering \citep{Lecavelier08b} seen at wavelengths shorter than $\sim$ 5,000, (2) a low temperature of $\pm$ 190 K at $\pm$ 0.27 mbar needed to explain the plateau at $\sim$ 1,500 km) above the minimum absorption level as a Na abundance change due to condensation, and finally (3) a high altitude $\sim$ 3,500 km) high temperature of $^{+500}_{-570}$ K at $\pm$ 0.003 mbar needed to fit the strong peak and narrow width of the cores of the Na I doublet. |
The pressure-teniperature gradient we fined when iuposiug condensation is close to and consistent (withiu lo) of an acliabatic eradieut. | The pressure-temperature gradient we find when imposing condensation is close to and consistent (within $\sigma$ ) of an adiabatic gradient. |
À temperature eradieut which is adiabatie or less is needed to be stable against convection. and solutious imposing adiabatic variations eive acceptable fits to the data (an acliabatic fit is shown in Fig. | A temperature gradient which is adiabatic or less is needed to be stable against convection, and solutions imposing adiabatic variations give acceptable fits to the data (an adiabatic fit is shown in Fig. |
3). | 3). |
We ultimately. however. choose to quote values found from the more general solution. as the luecar teniperature-atitude relation we asstune includes adiabatic eracdieuts axd the gveneral fit is located at a well hewaved ΜΗB Di. | We ultimately, however, choose to quote values found from the more general solution, as the linear temperature-altitude relation we assume includes adiabatic gradients and the general fit is located at a well behaved minimum $\chi^2$. |
Although the best fita solutionB isB «ligith: superacdiabatic. the relatively low quality of the claa still allows for a wide range of T-P eradieuts aud the aciabatic-like profile found from our ecucral fit indicates thi dan adiabatic behavior is likely at those altitudes. | Although the best fit solution is slightly superadiabatic, the relatively low quality of the data still allows for a wide range of T-P gradients and the adiabatic-like profile found from our general fit indicates that an adiabatic behavior is likely at those altitudes. |
Our model coutained two different amounts of atomic Na having separate variable parameters for he πακας ratios of the middle and lower atimospheres. | Our model contained two different amounts of atomic Na having separate variable parameters for the mixing ratios of the middle and lower atmospheres. |
The pressure boundaries (or condensation pressure) where tlic quantity changes occur is taken to be at an altitude where the temperature profile crosses the Naos condensation curve. | The pressure boundaries (or condensation pressure) where the quantity changes occur is taken to be at an altitude where the temperature profile crosses the $_2$ S condensation curve. |
Tn our fits. we found a Na musing? pane of 3.5!210)τον he middle atinosphere aud 241510© in the lower atinosphere. correspoudiug to 0.2 aud 2 nues solar abundance. | In our fits, we found a Na mixing ratio of $^{+2.9}_ {-1.9}\times10^{-7}$ in the middle atmosphere and $^{+2.1}_{-1.9}\times10^{-6}$ in the lower atmosphere, corresponding to 0.2 and 2 times solar abundance. |
Fortney(2005) estimated that the optical depth of Naos condensates iu transit geometry could be optically thick for IID209155b (7~0.73) for an atinosphere of solar composition aud cloud base at 30 mbar. | \cite{Fortney05} estimated that the optical depth of $_2$ S condensates in transit geometry could be optically thick for HD209458b $\tau\sim$ 0.73) for an atmosphere of solar composition and cloud base at 30 mbar. |
Given the fit T-P profile. we fud a Nas cloud 1vase of ~3 qub. | Given the fit T-P profile, we find a $_2$ S cloud base of $\sim$ 3 mbar. |
Lower pressures or larger particle sizes would provide a natural explanation for transparent clouds. as the condensate optical depth is proportional to the pressure and inversely proportional to the particle size (Alarley2000). | Lower pressures or larger particle sizes would provide a natural explanation for transparent clouds, as the condensate optical depth is proportional to the pressure and inversely proportional to the particle size \citep{Marley2000}. |
. Given our determined Na abundances. the mixing ratio of Χανς we find is (22-1) 1. which matches the assunied value of Fortuey(2005). | Given our determined Na abundances, the mixing ratio of $_2$ S we find is $\pm1$ $\times10^{-6}$, which matches the assumed value of \cite{Fortney05}. |
. Scaling the estimated optical depth of Fortnev(2005)... a cloud base at 3 nibar gives r90.073 . indicating that trauspareut Nas condensates are plausible. | Scaling the estimated optical depth of \cite{Fortney05}, a cloud base at 3 mbar gives $\tau\sim$ 0.073, indicating that transparent $_2$ S condensates are plausible. |
The coutiuuuuu πιονα of the Na D lines. between ~ L000 and 5.500Α.. shows a slightly lower absorption han the red-ward region. 7.000-8.000 À.. and a sienificant absorption “hump around ~6,250A wwhich appears in )oth resolutions. | The continuum blue-ward of the Na D lines, between $\sim$ 4,000 and 5,500, shows a slightly lower absorption than the red-ward region, 7,000-8,000, and a significant absorption `bump' around $\sim$ which appears in both resolutions. |
Our fits indicate hat the red-wrd alsorption levels (7.000-8,.000 Aj} are ikelv due to TiO iux VO at stratospheric heights. being arecly confined to atitudes below our Na condensation evel (also see Déósertetal. 2008)). | Our fits indicate that the red-ward absorption levels (7,000-8,000 ) are likely due to TiO and VO at stratospheric heights, being largely confined to altitudes below our Na condensation level (also see \citealt{desert08}) ). |
The TiO and VO catures mask part o the long wavoleugth Na line wine. hough not the shor waveleugth side. providiug strong constraints on the amudance and altitude of TiO aud VO. | The TiO and VO features mask part of the long wavelength Na line wing, though not the short wavelength side, providing strong constraints on the abundance and altitude of TiO and VO. |
The hot low atitude temperatures we find from Ravleigh scattering are consistent with the presence of TiO and VO. as boti species beein to condeuse out of the atinosphere at teuperatures below —1.500 EK. The low-altitude pressure assumed here is depeudeut on the detection of Ravleigh scattering (Lecavelicretal. 2008b).. thougyh unultiple interpretations of the absorption rise at those wavelengths exist (Lecavelierjuan 2007). | The hot low altitude temperatures we find from Rayleigh scattering are consistent with the presence of TiO and VO, as both species begin to condense out of the atmosphere at temperatures below $\sim$ 1,800 K. The low-altitude pressure assumed here is dependent on the detection of Rayleigh scattering \citep{Lecavelier08b}, though multiple interpretations of the absorption rise at those wavelengths exist \citep{Lecavelier08b,Ballester07,Bar07}. |
. Ultimaely. d£ quay be difficult without Ravleigh scattering to assign an absolute pressure scale to. these lueasurenieuts which are inherently differential. | Ultimately, it may be difficult without Rayleigh scattering to assign an absolute pressure scale to these measurements which are inherently differential. |
The effective planetary radius seen cing transit is scusitive to the quantity of the atimospheric pressure times the abundance of absorbing species (see equatious of Lecavelieretal. 2008a)). making an atmospheric fit at higher altitude possible by simply increasing the absorbers abundauuice. | The effective planetary radius seen during transit is sensitive to the quantity of the atmospheric pressure times the abundance of absorbing species (see equations of \citealt{Lecavelier08a}) ), making an atmospheric fit at higher altitude possible by simply increasing the absorbers abundance. |
The advantage of detecting Ravleigh scattering w IL when interpreting these models is that the abuvdance of Πω is known. constituting the bulk of the atinosphere. which thereby fixes the pressure. and defines the pressure scale for the rest of the trausinission spectrum. | The advantage of detecting Rayleigh scattering by $_2$ when interpreting these models is that the abundance of $_2$ is known, constituting the bulk of the atmosphere, which thereby fixes the pressure, and defines the pressure scale for the rest of the transmission spectrum. |
Our adoption of Ravleigh scattering here is based ou the goodness of fit to the data aud the necessity to probe a large pressure | Our adoption of Rayleigh scattering here is based on the goodness of fit to the data and the necessity to probe a large pressure |
bow shock distance. | bow shock distance. |
The two kinetic models display a plasma deceleration upstream of the bow shock due to charge exchange with component 3 and 4 neutrals that are streaming antisunward and have passed the bow shock. | The two kinetic models display a plasma deceleration upstream of the bow shock due to charge exchange with component 3 and 4 neutrals that are streaming antisunward and have passed the bow shock. |
To a lesser extent, the Flo and Mue models exhibit a similar deceleration, whereas the Sch model cannot resolve such counterstreaming fluid elements (they deposit their energy already far downstream of the BS). | To a lesser extent, the Flo and Mue models exhibit a similar deceleration, whereas the Sch model cannot resolve such counterstreaming fluid elements (they deposit their energy already far downstream of the BS). |
As a consequence of the deceleration both upstream and downstream of the BS, the BS is the weakest in the kinetic cases, followed in shock strength by the Flo model, and is the strongest in the one-fluid case, with the Mue model in between (Table 4)). | As a consequence of the deceleration both upstream and downstream of the BS, the BS is the weakest in the kinetic cases, followed in shock strength by the Flo model, and is the strongest in the one-fluid case, with the Mue model in between (Table \ref{resfull}) ). |
The shock-capturing method of BM arrives at a very weak BS. | The shock-capturing method of BM arrives at a very weak BS. |
This range of bow shock strengths explains the more gradual hydrogen wall in the kinetic cases. | This range of bow shock strengths explains the more gradual hydrogen wall in the kinetic cases. |
The hydrogen wall is of lesser amplitude in the kinetic cases because the velocities downstream of the BS are distinctly larger (absolute magnitude) than those in the fluid cases, and therefore charge-exchanged neutrals are not decelerated as much as in the one-fluid case where the plasma velocity is the smallest. | The hydrogen wall is of lesser amplitude in the kinetic cases because the velocities downstream of the BS are distinctly larger (absolute magnitude) than those in the fluid cases, and therefore charge-exchanged neutrals are not decelerated as much as in the one-fluid case where the plasma velocity is the smallest. |
To appreciate this difference, one has to mentally shift the plots of reffigvr (right) so that the individual bow shocks line up. | To appreciate this difference, one has to mentally shift the plots of \\ref{figvr} (right) so that the individual bow shocks line up. |
As expected, a stronger bow shock results in a more decelerated plasma, and therefore a larger peak density of the hydrogen wall. | As expected, a stronger bow shock results in a more decelerated plasma, and therefore a larger peak density of the hydrogen wall. |
Typically, this also means a lesser distance of the BS to the Sun, and this trend is evident in Table 4.. | Typically, this also means a lesser distance of the BS to the Sun, and this trend is evident in Table \ref{resfull}. |
The exception is model BM that experiences additional deceleration downstream of the BS, such that the BS standoff distance is not as far outward as the shock strength would suggest. | The exception is model BM that experiences additional deceleration downstream of the BS, such that the BS standoff distance is not as far outward as the shock strength would suggest. |
The different hydrogen walls result in different neutral densities downwind of the TS, where the neutrals enter region 4 of the supersonic solar wind. | The different hydrogen walls result in different neutral densities downwind of the TS, where the neutrals enter region 4 of the supersonic solar wind. |
The filtration ratio ng(TS)/ng(oo) is listed in Table 4,, along with the two absolute densities discussed above. | The filtration ratio $n_H(TS)/n_H(\infty)$ is listed in Table \ref{resfull}, along with the two absolute densities discussed above. |
In principle, it can be | In principle, it can be |
properties at such high redshifts was carried out by using both accretion and merger scenarios. | properties at such high redshifts was carried out by using both accretion and merger scenarios. |
They find that to produce sufficient high-mass SMBHs that power the most luminous SDSS quasars C».[0* MMo. their models yield à mass density of 10? MM MMpc? for SMBHs with masses around 105 MM. emphasising that this is too large by a factor of 100 to 1000. | They find that to produce sufficient high-mass SMBHs that power the most luminous SDSS quasars $\geq\,10^9$ $_{\sun}$ ), their models yield a mass density of $10^5$ $_{\sun}$ $^{-3}$ for SMBHs with masses around $10^8$ $_{\sun}$, emphasising that this is too large by a factor of 100 to 1000. |
Hence their estimate for such 105 MM; SMBHs is of the order 107 ΜΜΟ MMpe™. too. | Hence their estimate for such $10^8$ $_{\sun}$ SMBHs is of the order $10^2$ $_{\sun}$ $^{-3}$, too. |
We conclude that à non-negligible fraction of lower-mass SMBHs at the highest redshifts can put constraints on structure formation after the Big Bang — in particular considering SMBH growth models connected to the build-up of galaxies. | We conclude that a non-negligible fraction of lower-mass SMBHs at the highest redshifts can put constraints on structure formation after the Big Bang – in particular considering SMBH growth models connected to the build-up of galaxies. |
Aost of these theoretical frameworks focus only on the high-mass SMBHs necessary to power the most luminous quasars. | Most of these theoretical frameworks focus only on the high-mass SMBHs necessary to power the most luminous quasars. |
Present observations of lower-mass SMBHs are not possible using currently available instruments with reasonable integration times. but radio observations can play a crucial role to investigate the effects of black holes on various scales in the early universe. | Present observations of lower-mass SMBHs are not possible using currently available instruments with reasonable integration times, but radio observations can play a crucial role to investigate the effects of black holes on various scales in the early universe. |
We have investigated the implications of Infrared-Faint Radio Sources (IFRS) being high-redshift AGN. | We have investigated the implications of Infrared-Faint Radio Sources (IFRS) being high-redshift AGN. |
We draw consequences on AGN number counts and SMBH mass densities. the Cosmic X-ray Background (CXB) and theoretical models of structure formation after the Big Bang. | We draw consequences on AGN number counts and SMBH mass densities, the Cosmic X-ray Background (CXB) and theoretical models of structure formation after the Big Bang. |
Our main results are: Future observations of the IFRS population will further constrain the nature and thus the astrophysical significance of these objects. | Our main results are: Future observations of the IFRS population will further constrain the nature and thus the astrophysical significance of these objects. |
In particular sub-mm observatories such as ESO's ALMA array or the Space Observatory will be able to provide constraints on the MIR/FIR emission of IFRS. | In particular sub-mm observatories such as ESO's ALMA array or the Space Observatory will be able to provide constraints on the MIR/FIR emission of IFRS. |
They will clarify if the IFRS population predominantly consists of highly obscured type II AGN. or if IFRS in general show flatter SEDs without any distinct far-infrared bump produced by dust. | They will clarify if the IFRS population predominantly consists of highly obscured type II AGN, or if IFRS in general show flatter SEDs without any distinct far-infrared bump produced by dust. |
Corresponding deep observations have already been awarded time and will be executed in the forthcoming observational period. | Corresponding deep observations have already been awarded time and will be executed in the forthcoming observational period. |
They will put constraints on the star formation history of the universe by estimating dust masses for objects at redshifts potentially as high as 66. | They will put constraints on the star formation history of the universe by estimating dust masses for objects at redshifts potentially as high as 6. |
red- aud blue-shifted components combined. | red- and blue-shifted components combined. |
Ou Figs. | On Figs. |
2 and 3. (he arrows indicate where III nebular emission could have affected the profile. whereas on Fie. | 2 and 3, the arrows indicate where HI nebular emission could have affected the profile, whereas on Fig. |
4 the arrow indicates where (OL AG300 from the night skv may not have been removed completely. | 4 the arrow indicates where [OI] $\lambda$ 6300 from the night sky may not have been removed completely. |
Figure 2 shows that dramatic changes occurred in the Io emission line profile of IXII 15D during eclipse and egress. | Figure 2 shows that dramatic changes occurred in the $\alpha$ emission line profile of KH 15D during eclipse and egress. |
Out of eclipse. the star appears to be à WTTS with an EW(Ila ) οAL but with a double-peaked profile. having a central absorption and faint. but clearly detectable broad wings. | Out of eclipse, the star appears to be a WTTS with an $\alpha$ ) $\sim$ 2, but with a double-peaked profile, having a central absorption and faint, but clearly detectable broad wings. |
A comparison with the Ha profiles of other WTTSs shows that only 3 out of 19 such stars exhibited double-peaked profiles similar to KI 15D. with UX Tan A being the most similar (see Reipurthetal.(1996))). | A comparison with the $\alpha$ profiles of other WTTSs \citep{Hartmann82, Mundt83, FB87,
Edwards94, Reipurth96} shows that only 3 out of 19 such stars exhibited double-peaked profiles similar to KH 15D, with UX Tau A being the most similar (see \citet{Reipurth96}) ). |
Most WTTSs show narrow. single-peaked emission lines. | Most WTTSs show narrow, single-peaked emission lines. |
It is also interesting (hat the central absorption feature appears to extend below the stellar continuum. | It is also interesting that the central absorption feature appears to extend below the stellar continuum. |
This is unusual for auv TTS. weak ov classical. | This is unusual for any TTS, weak or classical. |
During eclipse and egress. (he "natural coronagraph effect is clear. | During eclipse and egress, the “natural coronagraph" effect is clear. |
Near the EW of Ha grows to ~40A. while the relative flux drops by ~50%.. | Near mid-eclipse, the EW of $\alpha$ grows to $\sim$ 40, while the relative flux drops by $\sim$. |
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