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For reference, the total ATLAS catalogue for this field contains ~ 6600 sources (Rigby et al. | For reference, the total ATLAS catalogue for this field contains $\sim$ 6600 sources (Rigby et al. |
2010, in prep.). | 2010, in prep.). |
The colour-colour plots are filled with 10° black-body spectra at a single dust temperature, 74, modified by a frequency-dependent emissivity function e,ος»?, where the flux density f, is In generating these models, we consider uniform ranges of dust temperature from 10 to 60 K, emissivity parameter 6 from 0 to 2, and redshift from0 to 5. | The colour-colour plots are filled with $^6$ black-body spectra at a single dust temperature, $T_{\rm d}$, modified by a frequency-dependent emissivity function $\epsilon_{\nu}\propto \nu^{\beta}$, where the flux density $f_\nu$ is In generating these models, we consider uniform ranges of dust temperature from 10 to 60 K, emissivity parameter $\beta$ from 0 to 2, and redshift from0 to 5. |
The choice of 0<82 compared to 1«82 makes a minor difference since we also broaden the SED tracks in the colour diagram by adding an extra Gaussian standard deviation of to the fluxes used to compute the model colour. | The choice of $0 < \beta <2$ compared to $1 < \beta < 2$ makes a minor difference since we also broaden the SED tracks in the colour diagram by adding an extra Gaussian standard deviation of to the fluxes used to compute the model colour. |
This scatter accounts for the broadening of data in the colour-colour plane caused by flux uncertainties. | This scatter accounts for the broadening of data in the colour-colour plane caused by flux uncertainties. |
As shown in Fig. | As shown in Fig. |
1, we find that the colour diagram of sources with fluxes only from SPIRE are well within the limits defined by the models we have considered. | 1, we find that the colour diagram of sources with fluxes only from SPIRE are well within the limits defined by the models we have considered. |
When we examine PACS colours, we find that some points lie outside the same set of tracks as used for the SPIRE-only colour diagram. | When we examine PACS colours, we find that some points lie outside the same set of tracks as used for the SPIRE-only colour diagram. |
While some of these outlier points may be cause by either the fractionally larger flux errors of PACS or contamination from a neighboring source, it is possible that some of these sources are not accurately described by our simple isothermal SED model, requiring for instance a second dust component (Dunne & Eales 2001) or a more complex SED model. | While some of these outlier points may be cause by either the fractionally larger flux errors of PACS or contamination from a neighboring source, it is possible that some of these sources are not accurately described by our simple isothermal SED model, requiring for instance a second dust component (Dunne $\&$ Eales 2001) or a more complex SED model. |
When fitting a simple modified black-body model to the data, we must keep in mind that there is a partial degeneracy between f and Τα and, more importantly, a perfect degeneracy between Τα and z. | When fitting a simple modified black-body model to the data, we must keep in mind that there is a partial degeneracy between $\beta$ and $T_{\rm d}$ and, more importantly, a perfect degeneracy between $T_{\rm d}$ and $z$. |
The peak of the SED is determined by the v/Ta term in the exponential, so that a measurement of the colours alone constrains only the ratio (1+z)/Ta. | The peak of the SED is determined by the $\nu/T_{\rm d}$ term in the exponential, so that a measurement of the colours alone constrains only the ratio $(1+z)/T_{\rm d}$. |
However, assuming reasonable priors on the free parameters of the SED model (in our case, 6 and Τα) it is still possible to estimate a qualitative redshift distribution for our sample of sources. | However, assuming reasonable priors on the free parameters of the SED model (in our case, $\beta$ and $T_{\rm d}$ ) it is still possible to estimate a qualitative redshift distribution for our sample of sources. |
Alternatively, if secure redshifts are known from optical identifications, we can determine the dust temperatures. | Alternatively, if secure redshifts are known from optical cross-identifications, we can determine the dust temperatures. |
We consider the H-ATLAS source sample with detections at 3c in at least two bands and at 5c in one band of either PACS and SPIRE that have been robustly (reliability parameter Εικ>0.9, Smith et al. | We consider the H-ATLAS source sample with detections at $3\sigma$ in at least two bands and at $5\sigma$ in one band of either PACS and SPIRE that have been robustly (reliability parameter $_{\rm LR}>0.9$, Smith et al. |
2010, in prep.) | 2010, in prep.) |
identified with GAMA or SDSS DR-7 galaxies. | identified with GAMA or SDSS DR-7 galaxies. |
We also require that there is a known spectroscopic redshift from either GAMA or SDSS, or from the photometric redshift catalog (with (1 z)/oc;>5 and ζ/σε> 1) that was generated for this field and cross-identified with ATLAS sources (Smith et al. | We also require that there is a known spectroscopic redshift from either GAMA or SDSS, or from the photometric redshift catalog (with $(1+z)/\sigma_z>5$ and $z/\sigma_z>1$ ) that was generated for this field and cross-identified with ATLAS sources (Smith et al. |
2010, in prep.). | 2010, in prep.). |
We select 330 sources, which correspond to the low redshift subsample of the galaxies selected in the colour-colour diagrams. | We select 330 sources, which correspond to the low redshift subsample of the galaxies selected in the colour-colour diagrams. |
For each galaxy, we perform a single temperature fit from the above equation assuming β= 1.5. | For each galaxy, we perform a single temperature fit from the above equation assuming $\beta=1.5$ . |
We found that an isothermal SED model generally is a good fit to the 330 galaxies. | We found that an isothermal SED model generally is a good fit to the 330 galaxies. |
Fitting a | Fitting a |
outside the virial radius. | outside the virial radius. |
When compared. to gas profiles inferred. from X-ray. data. our AGN model produces too shallow σας distribution. suggesting that we probably need even more powerful feedback. processes. | When compared to gas profiles inferred from X-ray data, our AGN model produces too shallow gas distribution, suggesting that we probably need even more powerful feedback processes. |
Our cluster formation simulations with AGN feedback have not fully converged vet - as we increase the resolution. we find a stellar mass profile for the BCC that is in better agreement with observations. but it is still too low by about a factor of two. | Our cluster formation simulations with AGN feedback have not fully converged yet - as we increase the resolution, we find a stellar mass profile for the BCG that is in better agreement with observations, but it is still too low by about a factor of two. |
We are still missing the lowest mass galaxy population. which could. provide the missing stellar mass in the central elliptical galaxy. | We are still missing the lowest mass galaxy population, which could provide the missing stellar mass in the central elliptical galaxy. |
We also note that. in the current picture. AGN feedback is a morphologically. dependent process: it only directly. alfects galaxies with SMDlIIs. ie. galaxies with a significant bulge/spheroid component. | We also note that, in the current picture, AGN feedback is a morphologically dependent process: it only directly affects galaxies with SMBHs, i.e. galaxies with a significant bulge/spheroid component. |
Higher resolution studies would be needed. in order to reliably model this clistinetion. so that star formation in diskv galaxies is not artificially suppressed. | Higher resolution studies would be needed, in order to reliably model this distinction, so that star formation in disky galaxies is not artificially suppressed. |
We thank our anonvmous referee for. helpful. suggestions that greatly. improved the quality of the paper. | We thank our anonymous referee for helpful suggestions that greatly improved the quality of the paper. |
ICE thanks Ancrey Ixravtsov for stimulating comments. | RT thanks Andrey Kravtsov for stimulating comments. |
All simulations were performed. on the Cray NT-5. cluster at CSCS. Alanno. Switzerland. | All simulations were performed on the Cray XT-5 cluster at CSCS, Manno, Switzerland. |
We thank the CES for supporting the astrophysical [Duids program at the University of Zürrich. | We thank the CTS for supporting the astrophysical fluids program at the University of Zürrich. |
1f one defines the initial radius of cach dark matter shell as r; and its final. aciabatically contracted. value ry. 2 haveproposed to capture ? model using the following simplified niocel The original can be recovered using a=1. | If one defines the initial radius of each dark matter shell as $r_i$ and its final, adiabatically contracted value $r_f$, \cite{Abadi:2009p531} haveproposed to capture \cite{Gnedin:2004p569} model using the following simplified model The original \cite{Blumenthal:1986p732} can be recovered using $\alpha =1$. |
The final cumulatec mass distribution7. is computed using where we have assumed that the initial dark matter mass ds conserved curing AC. | The final cumulated mass distribution is computed using where we have assumed that the initial dark matter mass is conserved during AC. |
The initial dark matter distribution is described using the analytical NEW profile where.—ης αμα ry=rsoof/6. | The initial dark matter distribution is described using the analytical NFW profile where $x=r_i/r_s$ and $r_s=r_{\rm 200}/c$. |
Mooo is the total virial mass. | $M_{\rm 200}$ is the total virial mass. |
For the barvonic cistribution. we assume a constant surface density cise with size ry and mass my. so that The dark matter mass fraction is computed using fy=Loma fAlooo. | For the baryonic distribution, we assume a constant surface density disc with size $r_d$ and mass $m_d$ , so that The dark matter mass fraction is computed using $f_d=1-m_d/M_{\rm 200}$ . |
The model we considered in Equation At for the barvonic mass distribution has been chosen that simple on purpose: inserting Equation A+ into the AC relation in Equation Al.. one clearly sees that we have to find the only real root of third order polynomial equation with unknown refi. | The model we considered in Equation \ref{equ:Baryons} for the baryonic mass distribution has been chosen that simple on purpose: inserting Equation \ref{equ:Baryons} into the AC relation in Equation \ref{equ:Abadi}, one clearly sees that we have to find the only real root of a third order polynomial equation with unknown $r_f/r_i$. |
This acan be done quite easily with any root finder. | This can be done quite easily with any root finder. |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one has | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one has | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$ |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_ |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \ |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \p |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \pr |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \pro |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \prop |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \propt |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \propto |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM; | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \propto |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM;x | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \propto x |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM;xa | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \propto x^ |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM;xa7 | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \propto x^2 |
In case the cise size is small enough. (namely ior,< ru). the AC model is fully tractable analvticallv by noticing that for we. one hasM;xa7. | In case the disc size is small enough (namely if $r_d \ll r_s$ ), the AC model is fully tractable analytically by noticing that for $x \ll 1$ , one has$M_i \propto x^2$ |
shows a weak and then strong peak in its light curve (Figure 1)), the initial peak powered by thermal energy and the second by the radioactive decay of ?9Ni. | shows a weak and then strong peak in its light curve (Figure \ref{restframeRlightcurve}) ), the initial peak powered by thermal energy and the second by the radioactive decay of $^{56}$ Ni. |
The heating from radioactive decay delays the inward-propagating recombination wave from ejecta cooling, regulating the electron scattering opacity (and thus the release of thermal energy), and causing the second peak to rise at 200-300 days, which reaches a spectacular brightness of ~—21.5 mag. | The heating from radioactive decay delays the inward-propagating recombination wave from ejecta cooling, regulating the electron scattering opacity (and thus the release of thermal energy), and causing the second peak to rise at 200-300 days, which reaches a spectacular brightness of $\sim -21.5$ mag. |
However, model R150 produces very little Ni, and therefore lacks a prominent second peak; the light curve is essentially thermally powered and reaches a brightness less than that of a Type Ia SN. | However, model R150 produces very little $^{56}$ Ni, and therefore lacks a prominent second peak; the light curve is essentially thermally powered and reaches a brightness less than that of a Type Ia SN. |
The helium core models are more compact and hence lack an initial thermal peak (Figure 2)). | The helium core models are more compact and hence lack an initial thermal peak (Figure \ref{restframeHelightcurve}) ). |
Model He130 reaches an exceptional peak brightness of 2x1043 ergs s!, whereas Model He80 demonstrates that despite being massive and energetic, not all PISNe are bright. | Model He130 reaches an exceptional peak brightness of $2\times 10^{44}$ ergs $^{-1}$, whereas Model He80 demonstrates that despite being massive and energetic, not all PISNe are bright. |
This steep mass-luminosity relation for PISNe suggests that to increase the sheer number of SNe detected, it is better to conduct a wide rather than deep survey of the sky (2).. | This steep mass-luminosity relation for PISNe suggests that to increase the sheer number of SNe detected, it is better to conduct a wide rather than deep survey of the sky \citep{Weinmann2005}. |
'The spectra of a PISN resemble that of average SNe, with P-Cygni line profiles on top of a blackbody, see Figures 3,, 4.. | The spectra of a PISN resemble that of average SNe, with P-Cygni line profiles on top of a blackbody, see Figures \ref{restframeR250}, \ref{restframeHe130}. |
For the RSG models, at early times, the spectrum is rather featureless with only weak Balmer and calcium lines, reflecting the low abundance of metals in unburned ejecta. | For the RSG models, at early times, the spectrum is rather featureless with only weak Balmer and calcium lines, reflecting the low abundance of metals in unburned ejecta. |
'The spectral energy distributions of the models are blue at earlier times (<50 days) but become redder over time as the expanding ejecta cools. | The spectral energy distributions of the models are blue at earlier times $\la 50$ days) but become redder over time as the expanding ejecta cools. |
In addition, line blanketing of the bluer wavelengths becomes more prominent over time, as the photosphere recedes into the deepest layers which are abundant in freshly synthesized iron group elements. | In addition, line blanketing of the bluer wavelengths becomes more prominent over time, as the photosphere recedes into the deepest layers which are abundant in freshly synthesized iron group elements. |
For PISNe at the redshifts of reionization, JWST will mostly be observing in the rest frame UV, so it is important to use more accurate spectral models, rather than the blackbody models of ?.. | For PISNe at the redshifts of reionization, JWST will mostly be observing in the rest frame UV, so it is important to use more accurate spectral models, rather than the blackbody models of \citet{Scannapieco2005a}. |
Spectroscopic or rest frame UV observations of PISNe may be able to constrain the metallicity of the progenitor star. | Spectroscopic or rest frame UV observations of PISNe may be able to constrain the metallicity of the progenitor star. |
However, the hydrogen envelope may be polluted by newly synthesized metals mixed out during the explosion. | However, the hydrogen envelope may be polluted by newly synthesized metals mixed out during the explosion. |
? simulated multi-dimensional models of PISNe to predict the degree of mixing. | \citet{Chen2011} simulated multi-dimensional models of PISNe to predict the degree of mixing. |
They found relatively small fluid instabilities generated fromburning at the boundaries of the oxygen shell, and concluded that PISNe keep their onion-shell structure in the explosion, until the reverse shock passes which generates Rayleigh-Taylor instabilities. | They found relatively small fluid instabilities generated fromburning at the boundaries of the oxygen shell, and concluded that PISNe keep their onion-shell structure in the explosion, until the reverse shock passes which generates Rayleigh-Taylor instabilities. |
This is in contrast with CCSNe, in which a shock runs through the inner metal-rich core, inducing the growth of instabilities and mixing. | This is in contrast with CCSNe, in which a shock runs through the inner metal-rich core, inducing the growth of instabilities and mixing. |
Also, ordinary Pop II/I CCSNe have non-zero metallicity in their hydrogen envelopes to begin with. | Also, ordinary Pop II/I CCSNe have non-zero metallicity in their hydrogen envelopes to begin with. |
Hence, metal lines in early-time spectroscopy might be able to distinguish PISNe from CCSNe, before the photosphere has receded deep into the ejecta. | Hence, metal lines in early-time spectroscopy might be able to distinguish PISNe from CCSNe, before the photosphere has receded deep into the ejecta. |
With a little mixing, N and possibly some C and O might appear in the early spectra of PISNe, but PISNe will not have any Si, Ni, Fe lines (?).. | With a little mixing, N and possibly some C and O might appear in the early spectra of PISNe, but PISNe will not have any Si, Ni, Fe lines \citep{Joggerst2011}. |
This best applies to the red supergiant models, as the Helium core models undergo significant burning and have spectra that show many metal lines at maximum light. | This best applies to the red supergiant models, as the Helium core models undergo significant burning and have spectra that show many metal lines at maximum light. |
Observations of quasar absorption spectra (?) indicate that reionization was completed by z= 6. | Observations of quasar absorption spectra \citep{Fan2006} indicate that reionization was completed by $z=6$ . |
It is believed | It is believed |
the galaxy. so considering more realistic galaxy shapes may be the key. | the galaxy, so considering more realistic galaxy shapes may be the key. |
While most loss cone studies have concentrated on static. spherical galaxy models. theory. and observations both indicate that dark matter halos and elliptical galaxies are at least mildly triaxial (Bak Statler 2000. Frans. Hlingworth. de Zeeuw 1991). | While most loss cone studies have concentrated on static, spherical galaxy models, theory and observations both indicate that dark matter halos and elliptical galaxies are at least mildly triaxial (Bak Statler 2000, Franx, Illingworth, de Zeeuw 1991). |
Triaxiality is present not only elliptical ealaxies. but in disc galaxies as well: a barred galaxy is a prime example ofa rotating triaxial ellipsoid. which has a more complex orbital structure aud response to a SMDII (Hasan Norman 1990. Sellwood Shen 2004). | Triaxiality is present not only elliptical galaxies, but in disc galaxies as well: a barred galaxy is a prime example of a rotating triaxial ellipsoid, which has a more complex orbital structure and response to a SMBH (Hasan Norman 1990, Sellwood Shen 2004). |
Over 10 percent of the local disc galaxy population is barred. and early indications suggest that this fraction may stav constant out to redshift 1 (Jogee οἱ al. | Over 70 percent of the local disc galaxy population is barred, and early indications suggest that this fraction may stay constant out to redshift 1 (Jogee et al. |
2004: Sheth et al. | 2004; Sheth et al. |
2004). | 2004). |
Our own galaxy hosts a bar with a semi-major axis length of about 2 kpc seen nearly end on [rom our perspective (O.jap~207: see Gerhard 2001 [or a review). | Our own galaxy hosts a bar with a semi-major axis length of about 2 kpc seen nearly end on from our perspective $\phi_{\odot -{\rm bar}} \sim 20^\circ$; see Gerhard 2001 for a review). |
Non-axisvmnmeltiry can introduce more stars to the loss cone in several wavs. | Non-axisymmetry can introduce more stars to the loss cone in several ways. |
First. stus in even a mildly (riaxial potential move in entirely different orbit families (han are present in a spheroid (see Fieure 1). | First, stars in even a mildly triaxial potential move in entirely different orbit families than are present in a spheroid (see Figure 1). |
In particular. there are a rich variety of regular box ancl boxlet orbits (hat are centrophilic and. comprise the backbone of the ealaxv (Miralda-Escude Schwarzschild 1989). | In particular, there are a rich variety of regular box and boxlet orbits that are centrophilic and comprise the backbone of the galaxy (Miralda-Escude Schwarzschild 1989). |
These centrophilic orbits can pass formally through. or very near {ο the SAIBIL which make them (he primary regular orbital component of the loss cone in a non axisvmmetric galaxy. | These centrophilic orbits can pass formally through, or very near to the SMBH, which make them the primary regular orbital component of the loss cone in a non axisymmetric galaxy. |
Besides the regular box ancl boxlet orbits. chaotic orbits often occur a (riaxial galaxy. | Besides the regular box and boxlet orbits, chaotic orbits often occur a triaxial galaxy. |
Though chaotic orbits are generically in a triaxial galaxy. {μον are particularly prevalent in one that is embedded with a SMDBIL | Though chaotic orbits are generically in a triaxial galaxy, they are particularly prevalent in one that is embedded with a SMBH. |
In fact. it is a commonly thought that SMDIIs cannot exist in a stable (riaxial galaxy. because black holes induce chaos in the centrophilic orbit families (Norman. May. van Albada 1985. Gerhard Binney 1985. Miralda-Escude Schwarzschild 1989. Merritt Quinlan 1993. Valluri Merritt. 1998. IHollev-Dockelmann el al. | In fact, it is a commonly thought that SMBHs cannot exist in a stable triaxial galaxy, because black holes induce chaos in the centrophilic orbit families (Norman, May van Albada 1985, Gerhard Binney 1985, Miralda-Escude Schwarzschild 1989, Merritt Quinlan 1998, Valluri Merritt 1998, Holley-Bockelmann et al. |
2002. Poon Merritt. 2002). | 2002, Poon Merritt 2002). |
and drive the galaxy toward: axisvinmetry in a [ew crossing limes. | and drive the galaxy toward axisymmetry in a few crossing times. |
Fully sell-consistent. high resolution n-body simulations show. however. that Óriaxialitv is not enlirely destroved by a SMDII. | Fully self-consistent, high resolution n-body simulations show, however, that triaxiality is not entirely destroyed by a SMBH. |
Though SMDIIS do indeed incite chaos in box and boxlet orbits at small radii and do rounden the region inside the radius οἱ inlluence riCM,o7. most of the centrophilic orbits are simply scattered onto resonant (or stickv) orbits (hat. allow (riaxialitv (o persist over most of the galaxy for many IIubble times (Ilollev-Bockelmann et al. | Though SMBHs do indeed incite chaos in box and boxlet orbits at small radii and do rounden the region inside the radius of influence $r_{\rm inf} \equiv G M_\bullet / \sigma^2 $, most of the centrophilic orbits are simply scattered onto resonant (or sticky) orbits that allow triaxiality to persist over most of the galaxy for many Hubble times (Holley-Bockelmann et al. |
2002: see also Poon Merritt 2002). | 2002; see also Poon Merritt 2002). |
Due to their stochastic motion through phase space. chaotic orbits can refill the loss cone as well. | Due to their stochastic motion through phase space, chaotic orbits can refill the loss cone as well. |
Stars on chaotic orbits encounter (he SAIBIT once per crossing πιο with a | Stars on chaotic orbits encounter the SMBH once per crossing time with a |
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