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8 shows the decomposition for the H8 + [Om] and Ha line profiles. | \ref{fig:sepair} shows the decomposition for the $\beta$ + ] and $\alpha$ line profiles. |
The micro-amplified spectrum is clearly a flat continuum in the HG + [Om] spectral region while there is some evidence that the core (and not the wings) of the Ha emission line is micro-amplified. | The micro-amplified spectrum is clearly a flat continuum in the $\beta$ + ] spectral region while there is some evidence that the core (and not the wings) of the $\alpha$ emission line is micro-amplified. |
The same kind of analysis can be done using the pair (C,AB). | The same kind of analysis can be done using the pair (C,AB). |
However the line profile differences are more subtle (cf. | However the line profile differences are more subtle (cf. |
Figs. | Figs. |
and 3)) so that the value of M is closer to the value of A (i.e. H closer to 1) in Eqs. | and \ref{fig:specuvir}) ) so that the value of $M$ is closer to the value of $A$ (i.e. $\mu$ closer to 1) in Eqs. |
4 and 5, making the extracted spectra noisier. | 4 and 5, making the extracted spectra noisier. |
Only the spectra obtained in 2000, which show the most conspicuous profile differences, are considered here. | Only the spectra obtained in 2000, which show the most conspicuous profile differences, are considered here. |
A de-magnification of the continuum explains the observations (Sect. 3.1)) | A de-magnification of the continuum explains the observations (Sect. \ref{sec:microe}) ) |
and the resulting Fy and Fy, are illustrated in Fig. 9.. | and the resulting $F_{M}$ and $F_{M\mu}$ are illustrated in Fig. \ref{fig:sepac}. |
They are roughly similar to those derived from the pair (D,AB) but the part of the emission profile which is micro-amplified is different. | They are roughly similar to those derived from the pair (D,AB) but the part of the emission profile which is micro-amplified is different. |
While only a small part of the core of the emission is seen in the Fy, profile derived from the (D,AB) pair, the red wing of the emission profile is also observed in Fy, computed from (C,AB). | While only a small part of the core of the emission is seen in the $F_{M\mu}$ profile derived from the (D,AB) pair, the red wing of the emission profile is also observed in $F_{M\mu}$ computed from (C,AB). |
This red wing and the blue emission peak are not seen in Fy, suggesting that the high-velocity component of the resonance line is micro-deamplified like the continuum. | This red wing and the blue emission peak are not seen in $F_{M}$, suggesting that the high-velocity component of the resonance line is micro-deamplified like the continuum. |
This is not unexpected since, for a given Einstein radius, demagnification regions with relatively smooth µ variations can extend on larger scales than amplification regions (e.g. Lewis and Ibata 2004)). | This is not unexpected since, for a given Einstein radius, demagnification regions with relatively smooth $\mu$ variations can extend on larger scales than amplification regions (e.g. Lewis and Ibata \cite{lew04}) ). |
The emission line core, which appears in both Εμμ and Fy, should originate, at least in part, from a region more extended thanthe high-velocity component. | The emission line core, which appears in both $F_{M\mu}$ and $F_{M}$, should originate, at least in part, from a region more extended thanthe high-velocity component. |
From the 2005 near-infrared spectra (the decomposition of which is not shown), we measure A = 0.74 for both the H6+[Om] and Ha regions. | From the 2005 near-infrared spectra (the decomposition of which is not shown), we measure $A$ = 0.74 for both the $\beta$ ] and $\alpha$ regions. |
The condition Fy,(M)2F. is verified at the wavelengths of H8 or Ha with M~ 0.87, while Fyuy(M)=F, is obtained at the wavelength of the [Om] lines with M~RIV 1.03. | The condition $F_{M\mu}(M) \geq F_c$ is verified at the wavelengths of $\beta$ or $\alpha$ with $M \simeq$ 0.87, while $F_{M\mu}(M) \geq F_c $ is obtained at the wavelength of the ] lines with $M \simeq$ 1.03. |
The latter value is comparable to the value of M derived from the UV-visible resonance lines. | The latter value is comparable to the value of $M$ derived from the UV-visible resonance lines. |
Although the narrow line region may be partially resolved (Chantry and Magain 2007)), the [Om] emission lines are expected to originate from a larger region and then less affected by microlensing, making M~ 1.03 a more plausible estimate. | Although the narrow line region may be partially resolved (Chantry and Magain \cite{cha07}) ), the ] emission lines are expected to originate from a larger region and then less affected by microlensing, making $M \simeq$ 1.03 a more plausible estimate. |
In this case, with uj=A/M~ 0.72 for the micro-amplification factor of the continuum, the Balmer emission lines do appear in both Fy, and Fy. | In this case, with $\mu = A/M \simeq$ 0.72 for the micro-amplification factor of the continuum, the Balmer emission lines do appear in both $F_{M\mu}$ and $F_{M}$. |
This means that they are also deamplified although not as much as the continuum. | This means that they are also micro-deamplified although not as much as the continuum. |
In principle, the variation of M with the wavelength can be attributed to differential extinction. | In principle, the variation of $M$ with the wavelength can be attributed to differential extinction. |
Unfortunately, for the (D,AB) pair, the wavelength dependence is not clear enough to extract an extinction curve, given the uncertainties on the determination of M (Fig. | Unfortunately, for the (D,AB) pair, the wavelength dependence is not clear enough to extract an extinction curve, given the uncertainties on the determination of $M$ (Fig. |
5 and Table 2)). | \ref{fig:ampli} and Table \ref{tab:ampli}) ). |
The results nevertheless suggest that the differential extinction between AB and D is lower than between A and B for which the extinction at Lya is~ 1.2 times the extinction at Ha (Fig. 4)). | The results nevertheless suggest that the differential extinction between AB and D is lower than between A and B for which the extinction at $\alpha$ is$\sim$ 1.2 times the extinction at $\alpha$ (Fig. \ref{fig:ratio}) ). |
As a consequence, the value M(D,AB) = 0.395+0.015 determined at thewavelength of Ho, i.e.in the reddest part of our spectra, should not differ from the true macro-amplification factor by more than 2%. | As a consequence, the value $M$ (D,AB) = $\pm$ 0.015 determined at thewavelength of $\alpha$ , i.e.in the reddest part of our spectra, should not differ from the true macro-amplification factor by more than . |
. For the (C,AB) pair, we conservatively adopt M(C,AB) = 0.95 + 0.08 at the wavelength of Ha. | For the (C,AB) pair, we conservatively adopt $M$ (C,AB) = 0.95 $\pm$ 0.08 at the wavelength of $\alpha$ . |
Comparing with M(C,AB) = 1.15 + 0.05 and M(C,AB) = 1.05 + 0.05 at the | Comparing with $M$ (C,AB) = 1.15 $\pm$ 0.05 and $M$ (C,AB) = 1.05 $\pm$ 0.05 at the |
but the relation σι,Alp ds steeper for non-Gaussian eroups. indicating a departure from the energy equipartition expectation. 0,x19OAL,S (soe FigureM 1). | but the relation $\sigma_u-M_R$ is steeper for non-Gaussian groups, indicating a departure from the energy equipartition expectation – $\sigma_u \propto 10^{0.2M_R}$ (see Figure 1). |
Our work suggests that the slope of the relation m,Alpe could be used to determine the evolutionary stage of galaxy. &roups. | Our work suggests that the slope of the relation $\sigma_u-M_R$ could be used to determine the evolutionary stage of galaxy groups. |
We thank the referee for verv useful. suggestions. | We thank the referee for very useful suggestions. |
We also thank S. Rembolcl for interesting discussions. | We also thank S. Rembold for interesting discussions. |
ALDI thanks the support of. CNPq. grants. 201322/2007-2. and 411254/2008-8. | ALBR thanks the support of CNPq, grants 201322/2007-2 and 471254/2008-8. |
PAAL thanks the support. of EAXPELBRJ. process 110.237/2010. | PAAL thanks the support of FAPERJ, process 110.237/2010. |
NIE thanks the support of FAPESP. process 2008/50198-2. | MT thanks the support of FAPESP, process 2008/50198-3. |
As seen in Table 6.. either entire atmosphere forecast or free atmosphere forecast has moderate error comparite with tje observation. while the former ends to slightly overestimate the value xl the latter ends to he opposition. | As seen in Table \ref{tbl-6}, either entire atmosphere forecast or free atmosphere forecast has moderate error comparing with the observation, while the former tends to slightly overestimate the value and the latter tends to the opposition. |
However. the tendeney is severe fo‘the cases such as Cerro machióun which the mea1 difference reads +0.167. | However, the tendency is severe for the cases such as Cerro Pachónn which the mean difference reads +0.46". |
A possibe explanation is that mauy observatories 'e located iu nountaltous areas and being actually muei higher than t1ο height of its belonged erid used in GES models (tlie height "di[erences vary rom 0-|«0001 auxng our sample. as shown LL Table 3)): Ilis may produce error i PBL seeing estilation. | A possible explanation is that many observatories are located in mountainous areas and being actually much higher than the height of its belonged grid used in GFS models (the height “differences” vary from 0-4,000m among our sample, as shown in Table \ref{tbl-3}) ); this may produce error in PBL seeing estimation. |
However. we find no substantial support for liis assuimpion as there is 110 siguilicantly edleney between the forecast bias and the height differenο (Figure 1)). | However, we find no substantial support for this assumption as there is no significantly tendency between the forecast bias and the height difference (Figure \ref{fig-1}) ). |
Draimaticaly. tlie site witl Siallest height clierence (Cerro Pachónn. with a height «lITereuce of 911) has the aregest nean errx (2-0.167). | Dramatically, the site with smallest height difference (Cerro Pachónn, with a height difference of 9m) has the largest mean error (+0.46"). |
Another ectally plausible explauation is the combiliec ellect of the poor consistency. between he GES/AXP inodel versus the actu:d CX4? variance. in PBL and error induced by the layer degeneration o ‘the ANP ιούς for hie1 altitude areas. | Another equally plausible explanation is the combined effect of the poor consistency between the GFS/AXP model versus the actual $C_{N}^{2}$ variance in PBL and error induced by the layer degeneration of the AXP model for high altitude areas. |
Iun fac. we do lind a weak bias tendeucy or [ree alinosprere forecast as shown i Figure 2:: the fo(απ for low-altitude sites tends to be jetter and a| 5ies below 3.000m have ai absolute mean forecas error below 0.257. | In fact, we do find a weak bias tendency for free atmosphere forecast as shown in Figure \ref{fig-2}: the forecast for low-altitude sites tends to be better and all sites below 3,000m have an absolute mean forecast error below 0.25". |
This feature fits e [act that tle al‘borne data used to determine the coel16161 sin the ANP model only iucludes ites with altitide ip to 2.5321n. But generally speaking. uo sbstantial correlation between bias πο ancl a llique geographic mocleling factor cau be ide1iliecd. therefore we can only suggest at the bias is conributed by a combination elfect of some LbiCOs. | This feature fits the fact that the airborne data used to determine the coefficients in the AXP model only includes sites with altitude up to 2,835m. But generally speaking, no substantial correlation between bias tendency and an unique geographic modeling factor can be identified, therefore we can only suggest that the bias is contributed by a combination effect of some factors. |
Το give a more comprehensive uucerstandiug ou the qualiM7 of the GES/ANP forecast. we plot e cumulative clist‘bution of the relative forecast error for each site as in Figure 3 (forecast for eutire atimospliere) aud Figure 1 {forecast for [ree atinosphie'e only). | To give a more comprehensive understanding on the quality of the GFS/AXP forecast, we plot the cumulative distribution of the relative forecast error for each site as in Figure \ref{fig-3} (forecast for entire atmosphere) and Figure \ref{fig-4} (forecast for free atmosphere only). |
Let ej to be the forecasted seeiug value ail eg to be the observed seeing value. the relative forecast error {ερ} is calculated by From the two figures.Om we can identify that the probabilities of producingOm forecast with <30% error concentrate in for entire atmosphere seeing forecast. aud for free atmosphere seeing forecast. | Let $\hat{\epsilon_{0}}$ to be the forecasted seeing value and $\epsilon_{0}$ to be the observed seeing value, the relative forecast error $E(\epsilon_{0})$ is calculated by From the two figures, we can identify that the probabilities of producing forecast with $<30$ error concentrate in for entire atmosphere seeing forecast and for free atmosphere seeing forecast. |
By contrast. Trinquet Vernin gives a probability of and for the o‘ginal AXP model to produce euire atinosphere or [ree atinosphere forecast in the same quality. | By contrast, Trinquet Vernin gives a probability of and for the original AXP model to produce entire atmosphere or free atmosphere forecast in the same quality. |
Geterally speaking. our result is iu ‘ough agreement with them. althoieh the models performance is rather uusatislyiug on a few part‘ular sites. | Generally speaking, our result is in rough agreement with them, although the model's performance is rather unsatisfying on a few particular sites. |
In addition to the suarizing table aud cumulative distribution figures. we also present the forecast-observation distriΠοια» for each site as iu Figure 5 aud Figure 6.. | In addition to the summarizing table and cumulative distribution figures, we also present the forecast-observation distributions for each site as in Figure \ref{fig-5} and Figure \ref{fig-6}. |
We can see tliat although the statistical vites αἱ be satisfying. the [ο'ecast-observation distribution figures ἱμαρίν that the correlation between forecast and observation Lis poor. | We can see that although the statistical values might be satisfying, the forecast-observation distribution figures imply that the correlation between forecast and observation is poor. |
This is not out-of-expected. since the study of Cherubinietal.(2009) at MINWC has sipgested that oue may achieve a good approximation of the actual condition with ~80 vertical lavers up to 10hPa. the ~15 layers used iu | This is not out-of-expected, since the study of \citet{che09} at MKWC has suggested that one may achieve a good approximation of the actual condition with $\sim80$ vertical layers up to 10hPa, the $\sim15$ layers used in |
and collaborators. | and collaborators. |
That is we consider (he (ransonie accretion of matter with the proper angular momentum (o produce a standing shock at a radius close to the horizon. which. subsequently accretes onto the black hole after passing through a downstream sonic point. | That is we consider the transonic accretion of matter with the proper angular momentum to produce a standing shock at a radius close to the horizon, which, subsequently accretes onto the black hole after passing through a downstream sonic point. |
Previous steady-state analysis found (hat itis possible to judiciously choose the specific angular momentum of the flow. so that the outer transonic one could be connected through a shock (ransition (o an inner (ransonic one which. passing through an inner sonic point. accretes onto the black hole. | Previous steady-state analysis found that it is possible to judiciously choose the specific angular momentum of the flow, so that the outer transonic one could be connected through a shock transition to an inner transonic one which, passing through an inner sonic point, accretes onto the black hole. |
The location of the shock in this situation is determined by finding a radial position at which the density and velocity of the two flow sections were those demanded by (he dissipative Rankine-Lligoniot conditions across a shock. | The location of the shock in this situation is determined by finding a radial position at which the density and velocity of the two flow sections were those demanded by the dissipative Rankine-Hugoniot conditions across a shock. |
So lar. the formation of standing shocks in hydrodynamic accretion has been extensively studied by a ΠΠο of authors for both inviscid flows (e.g..Chakrabarti1990:Sponholz&Molteni1994:Chakrabarti1996:Lu&Yuan1998:FukumuraTsuruta2004) and viscous flows (e.g...Chakrabarti1990:Luοἱal.1999:ChakrabarG&Das 2004).. considering either dissipative or non-dissipative shock jump conditions. | So far, the formation of standing shocks in hydrodynamic accretion has been extensively studied by a number of authors for both inviscid flows \citep[e.g.,][]{Cha90,Sponholz94,Cha96,Lu98,FT04} and viscous flows \citep[e.g.,][]{Cha90,Lu99,Cha04}, considering either dissipative or non-dissipative shock jump conditions. |
Recently standing shocks in the presence of poloidal magnetic fields [i.e.. magnetohvdrodynamic (MIID) shocks] around. a black hole was also studied for various parameter dependence (see Das Chakrabarti 2007 [ον Newtonian geometry: Takahashi et al. | Recently, standing shocks in the presence of poloidal magnetic fields [i.e., magnetohydrodynamic (MHD) shocks] around a black hole was also studied for various parameter dependence (see Das Chakrabarti 2007 for pseudo-Newtonian geometry; Takahashi et al. |
2002. Fukunira et al. | 2002, Fukumura et al. |
2007 [or Kerr geometry). | 2007 for Kerr geometry). |
In the context of the particle acceleration via the first-order Fermi mechanism across a shock front. the production of shock-accelerated. relativistic protons were discussed in spherical accretion (Protheroe&IXazanas1983:Ellison1986).. while other authors have explored the relativistie outflows in ADAF with shocks (Le&Becker2004.2005). | In the context of the particle acceleration via the first-order Fermi mechanism across a shock front, the production of shock-accelerated relativistic protons were discussed in spherical accretion \citep[][]{Kazanas83,Kazanas86}, while other authors have explored the relativistic outflows in ADAF with shocks \citep[][]{Becker04,Becker05}. |
. Similar altempts have been made to make plivsical connections between (he shocked-accretion ancl outflows. | Similar attempts have been made to make physical connections between the shocked-accretion and outflows. |
For instance. mass outflow rate were estimated [rom adiabatie shocked-flow region in Newlonian gravily (e.g..Chnakrabarti1999:Das2000). | For instance, mass outflow rate were estimated from adiabatic shocked-flow region in Newtonian gravity \citep[e.g.,][]{Cha99,Das00}. |
. Das&Chakrabarti(1999). took a similar approach to study the shock-generated. outflows with little energy. dissipation in pseudo-Newtonian geometry. | \citet{Das99} took a similar approach to study the shock-generated outflows with little energy dissipation in pseudo-Newtonian geometry. |
Independentlv. from general relativistic MIID simulations. the formation of jets (magneticall-drven and gas-pressure driven jets) is lound in the hieh pressure regions due to the shock/adiabatic compression (e.g..IXoideetal. 2005). | Independently, from general relativistic MHD simulations, the formation of jets (magnetically-driven and gas-pressure driven jets) is found in the high pressure regions due to the shock/adiabatic compression \citep[e.g.,][]{Koide99,Nishikawa05}. |
. They concluded that the jets are mainly produced by the gas-pressure gradient which is greatly. enhanced by the shock front at around ro6 gravitational radii (note that this feature was not seen in the Newtonian ealeulations). | They concluded that the jets are mainly produced by the gas-pressure gradient which is greatly enhanced by the shock front at around $r \sim 6$ gravitational radii (note that this feature was not seen in the Newtonian calculations). |
These studies also suggest an essential connection between (he shocked accretion flows and (he jets: i.e. the shock front may serve as a base of the outflows. | These studies also suggest an essential connection between the shocked accretion flows and the jets; i.e., the shock front may serve as a base of the outflows. |
The novelty of our approach lies in considering the possibili of shock formation (i.e. obeving the general relativistic. dissipative BRankine-IHIugoniot conditions at a shock οί) in which part of the mass. angular momentum and energy fluxes escape in the z-direction and do not participate in the shock transition. | The novelty of our approach lies in considering the possibility of shock formation (i.e. obeying the general relativistic, dissipative Rankine-Hugoniot conditions at a shock front) in which part of the mass, angular momentum and energy fluxes escape in the z-direction and do not participate in the shock transition. |
We (hen examine the energy per unit mass of the escaping matter which we compare to the escape velocity al the shock radius: if ilis greater (han the latter we conclude (hat this scenario can produce an outflow with its oulllow rate m and luminosity that are caleulable and can be compared to those of the entire accretion to obtain a measure of the elliciency of our “engine” in producing oulllows. | We then examine the energy per unit mass of the escaping matter which we compare to the escape velocity at the shock radius; if it is greater than the latter we conclude that this scenario can produce an outflow with its outflow rate $\dot m$ and luminosity that are calculable and can be compared to those of the entire accretion to obtain a measure of the efficiency of our “engine" in producing outflows. |
Some population of energetic nonthermal particles. produced via a shock acceleration. mav then be well separated from the equatorial accretion flows (which consists primarily of thermal particles). | Some population of energetic nonthermal particles, produced via a shock acceleration, may then be well separated from the equatorial accretion flows (which consists primarily of thermal particles). |
More specificallv. since we are interested in the formation of outflows through shocks in (he inner disk region relatively close to a presumably rotating black hole at a center (sav. rS30 gravitational radii). it is important to include the strong gravity and frame-drageing effects described by general relativity. | More specifically, since we are interested in the formation of outflows through shocks in the inner disk region relatively close to a presumably rotating black hole at a center (say, $r \lesssim 30$ gravitational radii), it is important to include the strong gravity and frame-dragging effects described by general relativity. |
To the best of our knowledge. no relevant work in | To the best of our knowledge, no relevant work in |
Supernova shock breakouts [rom normal massive stars (vpically eive rise to X-ray/ultraviolet bursts with a timescale <10* s | Supernova shock breakouts from normal massive stars typically give rise to X-ray/ultraviolet bursts with a timescale $\la 10^3$ s |
to other comets wilh extreme production-rate ratios. | to other comets with extreme production-rate ratios. |
The values in square brackets in Table 4 reler to logarithmic ratios between the production rates of various species. | The values in square brackets in Table \ref{Extreme} refer to logarithmic ratios between the production rates of various species. |
Where (here are discrepancies between the [N/OIL and the [N/CN] ratios for the comets. the [N/CN] ratios are more reliable. and are the basis for the listed [N/NII] ratios. | Where there are discrepancies between the [X/OH] and the [X/CN] ratios for the \citet{A+95} comets, the [X/CN] ratios are more reliable, and are the basis for the listed [X/NH] ratios. |
In 96P. we confirm the 907 report that the carbon-chain molecules are depleted compared to NIL | In 96P, we confirm the S07 report that the carbon-chain molecules are depleted compared to NH. |
Even compared to other "C-chain-depleted" comets 962 is unusual because we detect Cy emission but not CN. | Even compared to other “C-chain-depleted” comets 96P is unusual because we detect $_2$ emission but not CN. |
A distinctive property of 96P is the very high [Co/ CN] ratio. for which we only obtain a lower limit. | A distinctive property of 96P is the very high $_2$ /CN] ratio, for which we only obtain a lower limit. |
Rather than being carbon-chain depleted relative to CN. 96P is the comet with the least amount of CN to also display a C band. ancl by far the comet most depleted in CN relative to NIL. C5. and Cy than either "tvpical or "C-chain-depleted" comels. | Rather than being carbon-chain depleted relative to CN, 96P is the comet with the least amount of CN to also display a $_2$ band, and by far the comet most depleted in CN relative to NH, $_2$, and $_3$ than either “typical” or “C-chain-depleted” comets. |
Comet 96P/Machholz has extremely depleted Co and Cy relative to NIE and NIIS. with CN more extremely depleted. | Comet 96P/Machholz has extremely depleted $_2$ and $_3$ relative to NH and $_2$, with CN more extremely depleted. |
The dimensions of our spectrograph slit were sullicient (o see anv possible decay of the parents of these carbon molecules. | The dimensions of our spectrograph slit were sufficient to see any possible decay of the parents of these carbon molecules. |
The extreme carbon depletion of 96P is unlikelv to be due to cosmic rav. "gluüng in the comets past. and unlikely to be due to surface processing bv the Sun over repeated short orbits (since this effect is not seen in many other short-period comets). | The extreme carbon depletion of 96P is unlikely to be due to cosmic ray “gluing” in the comet's past, and unlikely to be due to surface processing by the Sun over repeated short orbits (since this effect is not seen in many other short-period comets). |
Though 96P was observed to have an outburst shortly alter its discovery. other confirmed outbursting comets show carbon emission leatures. | Though 96P was observed to have an outburst shortly after its discovery, other confirmed outbursting comets show carbon emission features. |
Thus. 1 appears (hat 96P/Machholz belongs to a small class of comets wilh genuinely. unusual molecule production rates. | Thus, it appears that 96P/Machholz belongs to a small class of comets with genuinely unusual molecule production rates. |
This research is supported by NASA erant NNGO5GG59G through the Planetary Astronomy progran. | This research is supported by NASA grant NNG05GG59G through the Planetary Astronomy program. |
temperature and surface metallicity, since such a procedure is closer to observational comparisons. | temperature and surface metallicity, since such a procedure is closer to observational comparisons. |
Two additional non-rotating stars of MMpo are then computed in order to obtain models with the same location in the HR diagram and surface metallicity as the rotating ΜΜο model studied before. | Two additional non-rotating stars of $_{\odot}$ are then computed in order to obtain models with the same location in the HR diagram and surface metallicity as the rotating $_{\odot}$ model studied before. |
For this purpose, the initial chemical composition and mixing-length parameter of non-rotating models are calibrated to reproduce simultaneously the surface metallicity and location in the HR diagram of the rotating models with an age of 3 and GGyr. | For this purpose, the initial chemical composition and mixing-length parameter of non-rotating models are calibrated to reproduce simultaneously the surface metallicity and location in the HR diagram of the rotating models with an age of 3 and Gyr. |
The evolutionary tracks of these models are shown in Fig. 11.. | The evolutionary tracks of these models are shown in Fig. \ref{dhr_ystrot}. |
The asteroseismic properties of rotating and non-rotating models are first compared by computing the mean large separation. | The asteroseismic properties of rotating and non-rotating models are first compared by computing the mean large separation. |
The rotating model exhibits a mean large separation of 150.1 and uHz after 3 and GGyr, while the corresponding non-rotating models have a mean large separation of 149.9 and Hz, respectively. | The rotating model exhibits a mean large separation of 150.1 and $\mu$ Hz after 3 and Gyr, while the corresponding non-rotating models have a mean large separation of 149.9 and $\mu$ Hz, respectively. |
These values are almost identical for rotating and non-rotating models since the mean large separation is mainly sensitive to the mean density of the star. | These values are almost identical for rotating and non-rotating models since the mean large separation is mainly sensitive to the mean density of the star. |
For this comparison, the models have the same mass and radius; they exhibit therefore a similar value of the mean large separation. | For this comparison, the models have the same mass and radius; they exhibit therefore a similar value of the mean large separation. |
Concerning the small separation, and more precisely the ratio ro» defined above, the situation is quite different. | Concerning the small separation, and more precisely the ratio $r_{02}$ defined above, the situation is quite different. |
Figure indeed shows an increase of the mean value of ro» for rotating models compared to non-rotating models with the same location in the HR diagram and surface metallicity. | Figure \ref{r02_t3t6} indeed shows an increase of the mean value of $r_{02}$ for rotating models compared to non-rotating models with the same location in the HR diagram and surface metallicity. |
As observed in Sect. | As observed in Sect. |
3.2.1 in models with the same initial parameters, this difference increases during the evolution of the star. | \ref{sec_modini} in models with the same initial parameters, this difference increases during the evolution of the star. |
The rotating model exhibits indeed a mean value of the ratio 7o» of 0.086 and 0.069 after 3 and GGyr, while the corresponding non-rotating models have a mean ratio of 0.083 | The rotating model exhibits indeed a mean value of the ratio $r_{02}$ of 0.086 and 0.069 after 3 and Gyr, while the corresponding non-rotating models have a mean ratio of 0.083 |
of parameters. most of the standard HΊσοςao discoveΝ chanels — those involving //—WW* aud hi—ZZ*. as well as all weak-bosou-fusion processes not ivolving direct Higgs decays to leptous — are strongly suppressed: furthermore. other contributiimο channels such as gg—hi55 and ΠΜ—bb) are also moderately suppressed. | of parameters, most of the standard Higgs discovery channels — those involving $h\rightarrow WW^\ast$ and $h\rightarrow ZZ^\ast$, as well as all weak-boson-fusion processes not involving direct Higgs decays to leptons — are strongly suppressed; furthermore, other contributing channels such as $gg\rightarrow h\rightarrow \gamma\gamma$ and $t\bar{t}h(h\rightarrow b\bar{b})$ are also moderately suppressed. |
La such a eise. the leptonic channels discussed iu Section 5 — especially ones such as ///i(//—77). which do uot involve a direct coupling between /i aud the electroweak gauge bosous — may well cousitute the only observable evidence of the Higgs boson. aud would thus be crucial for its discovery at the LHC. | In such a case, the leptonic channels discussed in Section \ref{sec:LHCSignatures} — especially ones such as $tth(h\rightarrow\tau\tau)$, which do not involve a direct coupling between $h$ and the electroweak gauge bosons — may well constitute the only observable evidence of the Higgs boson, and would thus be crucial for its discovery at the LHC. |
The phenomenology of a light HiOOees boson in Two-Hieesao-Doublet Models can cli[Ier drastically [rom that of a SM Higgs. | The phenomenology of a light Higgs boson in Two-Higgs-Doublet Models can differ drastically from that of a SM Higgs. |
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