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llowever theSpifzer images show much more details of (he outflow than the nenmr-infrared images.
However the images show much more details of the outflow than the near-infrared images.
In Figs.
In Figs.
5. and 6 we can see extended diffuse emission to the northeast ol EGO 34c and to the southwest of EGO 34a.
\ref{fig5} and \ref{fig6} we can see extended diffuse emission to the northeast of EGO 34c and to the southwest of EGO 34a.
Thecentral source of the outflow. MMS126. which is identified as a low-mass Class 0 object in the millimetre continuum observations bv Stankeοἱal.(2006). (Classified as Class0 in Evansetal. 2009)) is visible in all IRAC bands and also in the MIPS 24 iimage.
Thecentral source of the outflow, MMS126, which is identified as a low-mass Class 0 object in the millimetre continuum observations by \citet{sta06} (Classified as Class 0 in \citealt{eva09}) ) is visible in all IRAC bands and also in the MIPS 24 image.
Stankeοἱal.(2006) has observed a molecular CO outflow from MMS126 which is orientated in the NE-SW clirection.
\citet{sta06} has observed a molecular CO outflow from MMS126 which is orientated in the NE-SW direction.
Fies.
Figs.
20 and 21 present the images of EGOs 02. 06. 08. 15. and 16 that. corresponds to near-infrared flows [NGS2004) [10-01 and. [INGS2004| 10-02. (INhanzadyanetal. 2004).
\ref{fig20} and \ref{fig21} present the images of EGOs 02, 06, 08, 15, and 16 that corresponds to near-infrared flows [KGS2004] f10-01 and [KGS2004] 10-02 \citep{kha04}. .
. IXhanzadyanetal.(2004) detected nine Il» emission components of [INGS2004|. [10-01
\citet{kha04} detected nine $_{2}$ emission components of [KGS2004] f10-01.
surface density, more of the Πο sits at high densities in a low Z, high UV field, galaxy compared to a high Z, low UV field galaxy.
surface density, more of the $\H2$ sits at high densities in a low $Z$, high UV field, galaxy compared to a high $Z$, low UV field galaxy.
Furthermore, in the regime in which ng>Τις and the hydrogen gas is (close to) fully molecular the SFRs scale as οςnjj?.
Furthermore, in the regime in which $n_{\rm H}>\nc$ and the hydrogen gas is (close to) fully molecular the SFRs scale as $\propto{}n_{\rm H}^{1.5}$.
This then amounts to a higher SFR (and hence intercept) in a low Z, high UV field, galaxy compared to a high Z, low UV field galaxy.
This then amounts to a higher SFR (and hence intercept) in a low $Z$, high UV field, galaxy compared to a high $Z$, low UV field galaxy.
In other words, the large-scale SFRs depend not only on the large-scale H2 surface densities, but also on the distribution function of ny on small scales.
In other words, the large-scale SFRs depend not only on the large-scale $\H2$ surface densities, but also on the distribution function of $n_{\rm H}$ on small scales.
A related mechanism leads to a dependence of the slope on Z and Uww (left panel in the top and middle rows of Fig.2)).
A related mechanism leads to a dependence of the slope on $Z$ and $\UMW$ (left panel in the top and middle rows of \ref{fig:regfit}) ).
The panels in the top row of Fig.
The panels in the top row of Fig.
3 show that the typical hydrogen densities of cells that contribute to a given Hg surface density Uy, increase with Xg,.
\ref{fig:SH2NB} show that the typical hydrogen densities of cells that contribute to a given $\H2$ surface density $\Sigma_\H2$ increase with $\Sigma_\H2$.
This demonstrates that the density structure of disks changes with metallicity and radiation field of the ISM.
This demonstrates that the density structure of disks changes with metallicity and radiation field of the ISM.
The bottom panels in Fig.
The bottom panels in Fig.
3 show that this trend remains (although somewhat weakened) if the hydrogen density distribution is weighted by Hz mass.
\ref{fig:SH2NB} show that this trend remains (although somewhat weakened) if the hydrogen density distribution is weighted by $\H2$ mass.
It also shows, that the effect of an increasing hydrogen density with large-scale Πο surface density is stronger for a high Z, low Uww galaxy than for a low Z, high Umw galaxy.
It also shows, that the effect of an increasing hydrogen density with large-scale $\H2$ surface density is stronger for a high $Z$, low $\UMW$ galaxy than for a low $Z$, high $\UMW$ galaxy.
However, in the former case, most of the Hg is in cells with densities ng«n. and consequently the SFR density is still proportional to the total H2 density, i.e., slope 1.
However, in the former case, most of the $\H2$ is in cells with densities $n_{\rm H} < \nc$ and consequently the SFR density is still proportional to the total $\H2$ density, i.e., slope 1.
In the latter case, however, this increase is important.
In the latter case, however, this increase is important.
Let us see why.
Let us see why.
From (1) and (2)) it is clear that the surface density of the SFR is proportional to >);ng,ΤΗ, where the sum is over all cells within the given line-of-sight cylinder and a is either 0 (if ng«nc.) or 0.5 (if ng>nc).
From \ref{eq:SFR}) ) and \ref{eq:tau}) ) it is clear that the surface density of the SFR is proportional to $\sum_i n_{\H2,i} n_{{\rm H},i}^{\alpha_i}$ , where the sum is over all cells within the given line-of-sight cylinder and $\alpha$ is either 0 (if $n_{\rm H}<\nc$ ) or 0.5 (if $n_{\rm H}\geq{}\nc$ ).
The Ha surface density is proportional to $7;ng,,;.
The $\H2$ surface density is proportional to $\sum_i n_{\H2,i}$.
An increase in the H5 surface density Uy,—Xg, can be achieved in several ways.
An increase in the $\H2$ surface density $\Sigma_{\H2}\rightarrow{}\gamma{}\Sigma_{\H2}$ can be achieved in several ways.
If the length of the cylinder increases, then the surface density of the SFR increases proportional to ΣΗ..
If the length of the cylinder increases, then the surface density of the SFR increases proportional to $\Sigma_{\H2}$.
If, however, the density structure changes (in the simplest case via ny—"yng), then the surface density of the SFR increases by 4!** (assuming the gas is fully molecular).
If, however, the density structure changes (in the simplest case via $n_{\rm H}\rightarrow{}\gamma{}n_{\rm H}$ ), then the surface density of the SFR increases by $\gamma{}^{1+\alpha}$ (assuming the gas is fully molecular).
Obviously, if a=0 (as for ng«n.) the predicted large-scale slope of the relation is, as expected, linear.
Obviously, if $\alpha=0$ (as for $n_{\rm H}<\nc$ ) the predicted large-scale slope of the relation is, as expected, linear.
However, it lies between 1 and 1+a if there is a mix of cells with densities below and above nc.
However, it lies between 1 and $1+\alpha$ if there is a mix of cells with densities below and above $\nc$.
In addition, if the density distribution changes in a more complicated manner with Xg,, it is also possible to obtain "rolling" slopes or even large-scale slopes that are steeper than 1.5.
In addition, if the density distribution changes in a more complicated manner with $\Sigma_{\H2}$, it is also possible to obtain “rolling” slopes or even large-scale slopes that are steeper than 1.5.
We conclude thatdensity, see also Kravtsov(2003)..
We conclude that, see also \cite{2003ApJ...590L...1K}.
While the time averaging of the SFRs generates most of the scatter in the relation, the trends of slope and scatter with Z and— Uyw are largely driven by the non-linear coupling between SFR and Πο density.
While the time averaging of the SFRs generates most of the scatter in the relation, the trends of slope and scatter with $Z$ and $\UMW$ are largely driven by the non-linear coupling between SFR and $\H2$ density.
This can be clearly seen in the last row of Fig. 2..
This can be clearly seen in the last row of Fig. \ref{fig:regfit}.
If n,=109 cm-?, the slope of the relation changes only between 1.03 (7/79=1, Umw= 0.1) and 1.16 (Z/Ze= 0.1, Umw= 100) and the scatter increases only from 0.09 dex to 0.12 dex.
If $\nc=10^6$ $^{-3}$, the slope of the relation changes only between 1.03 $Z/Z_\odot=1$, $\UMW=0.1$ ) and 1.16 $Z/Z_\odot=0.1$ , $\UMW=100$ ) and the scatter increases only from 0.09 dex to 0.12 dex.
We discuss the dependence of the scatter on ISM properties further in the next section.
We discuss the dependence of the scatter on ISM properties further in the next section.
If SFRs are measured instantaneously the small-scale relation between star formation rate density and H» density is linear (i.e., nc is large), then the slope reduces to exactly unity, and any dependence of the intercept on metallicity or radiation field is eliminated and the scatter vanishes (at least as long as there are no other sources of scatter, see refsect:scale)).
If SFRs are measured instantaneously the small-scale relation between star formation rate density and $\H2$ density is linear (i.e., $\nc$ is large), then the slope reduces to exactly unity, and any dependence of the intercept on metallicity or radiation field is eliminated and the scatter vanishes (at least as long as there are no other sources of scatter, see \\ref{sect:scale}) ).
The reason lies in the fact that spatial averaging (which is a linear operation) over a linear relation between SFR and Ἡο density on small scales results again in a linear relation on larger scales.
The reason lies in the fact that spatial averaging (which is a linear operation) over a linear relation between SFR and $\H2$ density on small scales results again in a linear relation on larger scales.
We conclude that slope, intercept, and scatter of the relation averaged on kpc scales can change systematicallyUy, with metallicity and radiation field.
We conclude that slope, intercept, and scatter of the relation averaged on kpc scales can change systematically with metallicity and radiation field.
While our quantitative predictions likely depend on the assumedstar formation and the model for Hz formation and shielding the of a non-universality” of the large-scale relation is a rather generic outcome whenever there is a non-linear relation between SFR and H2 density on small scales.
While our quantitative predictions likely depend on the assumedstar formation and the model for $\H2$ formation and shielding the of a ``non-universality'' of the large-scale relation is a rather generic outcome whenever there is a non-linear relation between SFR and $\H2$ density on small scales.
svuchrotron seed. photons produced by the jet. plasma itself 2000).
synchrotron seed photons produced by the jet plasma itself \citep[e.g., see][]{bla00}.
In the rest frame of the emitting plasma. radiation from the torus and DLR depends on ocalion of the emission region in the jet and its bulk Lorentz [nctor I.
In the rest frame of the emitting plasma, radiation from the torus and BLR depends on location of the emission region in the jet and its bulk Lorentz factor $\Gamma'_p$.
Racliation from the torus is in general expected to be Doppler boosted in the plasma rest frame.
Radiation from the torus is in general expected to be Doppler boosted in the plasma rest frame.
This results in stronger dependence of the ERC emission on D, compared with the SSC enission. whose seed radiation is produced in the rest frame of the plasma.
This results in stronger dependence of the ERC emission on $\Gamma'_p$ compared with the SSC emission, whose seed radiation is produced in the rest frame of the plasma.
However. this behavior can be nodified by the effects of electron energy. losses from ERC radiation which. if dominant. imit the spatial extent of the emitüng plasma in an energv-dependent manner. leading to requency stratification of the emission.
However, this behavior can be modified by the effects of electron energy losses from ERC radiation which, if dominant, limit the spatial extent of the emitting plasma in an energy-dependent manner, leading to frequency stratification of the emission.
In (his paper. we investigate the broadband features of the ERC emission during flares under different assumptions about 1) properties of the molecular torus. 2) location in the jet where the flare occurs. and 3) the value of I, of the emitting plasma.
In this paper, we investigate the broadband features of the ERC emission during flares under different assumptions about 1) properties of the molecular torus, 2) location in the jet where the flare occurs, and 3) the value of $\Gamma'_p$ of the emitting plasma.
As in Paper I. we concentrate on the study of relative delavs between flares αἱ different [requencies and the features that can be used to distinguish among different emission mechanisms.
As in Paper I, we concentrate on the study of relative delays between flares at different frequencies and the features that can be used to distinguish among different emission mechanisms.
An additional eoal of this study is to define the properties of (he molecular torus and the BLR under which the results obtained in Paper I are valid.
An additional goal of this study is to define the properties of the molecular torus and the BLR under which the results obtained in Paper I are valid.
There we assumed that the ERC emission provides a neelieible contribution compared with the SSC radiation at the frequencies of interest. and that the energy losses of electrons are dominated by svnchrotron emission.
There we assumed that the ERC emission provides a negligible contribution compared with the SSC radiation at the frequencies of interest, and that the energy losses of electrons are dominated by synchrotron emission.
The inclusion of ERC radiation in our model of variability introduces a new component ol high-enerev emission.
The inclusion of ERC radiation in our model of variability introduces a new component of high-energy emission.
It ean also change the overall structure of the emission region and. hence. affect svnchrotron and SSC radiation if the energy losses of electrons are dominated bv scattering of external photons.
It can also change the overall structure of the emission region and, hence, affect synchrotron and SSC radiation if the energy losses of electrons are dominated by scattering of external photons.
In (his case Che decav time of electrons will be reduced compared wilh the case of pure svuchrotvon losses. which leads to changes in flux levels and values of critical frequencies. such as the break frequency.
In this case the decay time of electrons will be reduced compared with the case of pure synchrotron losses, which leads to changes in flux levels and values of critical frequencies, such as the break frequency.
The details depend on the structure and properties of (he sources of external emission. which we describe in (his section.
The details depend on the structure and properties of the sources of external emission, which we describe in this section.
As in Paper 1. here we concentrate on the study of rapid variability on time scales e»I]day.
As in Paper I, here we concentrate on the study of rapid variability on time scales $\sim1\,\mbox{day}$.
We adopt the assumptions about geometry and excitation structure of the emitting volume made in Paper 1. We ignore the expansion of the emission region. which is assumed to be a evlincder oriented along the jet.
We adopt the assumptions about geometry and excitation structure of the emitting volume made in Paper I. We ignore the expansion of the emission region, which is assumed to be a cylinder oriented along the jet.
We assume that the size of the emitting volume is small compared to the sizes of and distances to the external sources of emission.
We assume that the size of the emitting volume is small compared to the sizes of and distances to the external sources of emission.
Then we can assume (hat the external radiation is homogeneous throughout the emitting plasma.
Then we can assume that the external radiation is homogeneous throughout the emitting plasma,
Blueward of the Lyman limit lies a wavelength region of great observational uncertainty and astrophivsical significance.
Blueward of the Lyman limit lies a wavelength region of great observational uncertainty and astrophysical significance.
These Fu-UV photons control ionization of the interstellar and intergalactic media today. and at all times since recombination(Macauetal. 1999)..
These far-UV photons control ionization of the interstellar and intergalactic media today, and at all times since recombination\citep{mad99}. .
discussions.
discussions.
This work is supported in part by the Ministry of Education. Culture. Sports. Science. nn (MENT) Research AM Start-up 2284007.... (NIX) and Young Scientist (B3) 10115 (NO).
This work is supported in part by the Ministry of Education, Culture, Sports, Science, and Technology (MEXT) Research Activity Start-up 2284007 (NK) and Young Scientist (B) 20740115 (KO).
down to the infinitely small scale, to avoid a divergence of the total mass, the quantity nf?f, must be less than unity.
down to the infinitely small scale, to avoid a divergence of the total mass, the quantity $n f_r^3 f_{\rho}$ must be less than unity.
Similarly, for a halo with cut-off radius rmaz, halo to subhalo distance D must be greater than or equal to (1—fr)rmax/(12f,) to avoid any overlap.
Similarly, for a halo with cut-off radius $r_{max}$, halo to subhalo distance $D$ must be greater than or equal to $(1-f_r)r_{max}/(1-2f_r)$ to avoid any overlap.
Also note that the density of any subhalo at the cutoff radius P(rmas) scales as
Also note that the density of any subhalo at the cutoff radius $\rho(r_{max})$ scales as $f_{\rho}$.
So, we adopt f,=1.0 for further analysis to ensure a fj.constant density at the cutoff radius for all the substructures.
So, we adopt $f_{\rho} = 1.0$ for further analysis to ensure a constant density at the cutoff radius for all the substructures.
In this simplified model, we have assumed that all the parameters like D, 0o, Tmax etc.
In this simplified model, we have assumed that all the parameters like $D$, $\rho_0$, $r_{max}$ etc.
are identical for all the subhalos for a particular substructure level and n, f, and f, remain constant for all the levels.
are identical for all the subhalos for a particular substructure level and $n$, $f_r$ and $f_\rho$ remain constant for all the levels.
In a realistic situation, however, all these parameters may have some random variation.
In a realistic situation, however, all these parameters may have some random variation.
As a result of these fluctuations, the halo mass function, the radial mass distribution, and in turn, the rotation curves, are expected to be somewhat smoother than that are predicted from this analysis.
As a result of these fluctuations, the halo mass function, the radial mass distribution, and in turn, the rotation curves, are expected to be somewhat smoother than that are predicted from this analysis.
A much stronger constraint on the parameters comes from the consideration of the stability of this structures preventing tidal disruption by invoking self-gravity.
A much stronger constraint on the parameters comes from the consideration of the stability of this structures preventing tidal disruption by invoking self-gravity.
Considering a rigid object of mass m and radius r at à distance d from a bigger object of mass M and radius R, from the standard Roche limit consideration, D should be ~r(2M/m)!/? to avoid tidal disruption of the smaller body.
Considering a rigid object of mass $m$ and radius $r$ at a distance $d$ from a bigger object of mass $M$ and radius $R$, from the standard Roche limit consideration, $D$ should be $\sim r(2M/m)^{1/3}$ to avoid tidal disruption of the smaller body.
In the case of any halo and its immediate subhalo, m/M=f? and r/R=f. imply d©1.25R.
In the case of any halo and its immediate subhalo, $m/M = f_r^3$ and $r/R = f_r$ imply $d \approx 1.25 R$.
For even smaller subhalo structures with r/R=fF, the mass is scaled accordingly m/M=f?* and de1.25R assures stability.
For even smaller subhalo structures with $r/R = f_r^k$, the mass is scaled accordingly $m/M = f_r^{3k}$ and $d \approx 1.25 R$ assures stability.
Hence, a distance will make the whole structure stable.
Hence, a distance will make the whole structure stable.
Here x is a fudge factor and we use the value of x=1.1 to accommodate non-rigid density clumps.
Here $x$ is a fudge factor and we use the value of $x=1.1$ to accommodate non-rigid density clumps.
For a given value of D, the number of substructure n is also constrained to be to avoid any possible overlap of subhalo structures.
For a given value of $D$, the number of substructure $n$ is also constrained to be to avoid any possible overlap of subhalo structures.
A detailed virial stability analysis requires numerical simulation of the dynamics of such a density distribution to get the time-averaged dynamical quantities.
A detailed virial stability analysis requires numerical simulation of the dynamics of such a density distribution to get the time-averaged dynamical quantities.
But, à simple ensemble average virial scaling analysis may be used to constrain the central density po for a set of model parameter.
But, a simple ensemble average virial scaling analysis may be used to constrain the central density $\rho_0$ for a set of model parameter.
As the whole structure is assumed to have an approximate spherical symmetry, average kinetic and potential energies, (T) and (V), for thin spherical shells of radius r and thickness ór will bewhere opm is dark matter velocity dispersion, M(r) is total mass within radius r and = is the scale dependent virial velocity v.(r)equivalent (GM/r)!?of the "circular velocity" for rotating disk.
As the whole structure is assumed to have an approximate spherical symmetry, average kinetic and potential energies, $\langle T \rangle$ and $\langle V \rangle$, for thin spherical shells of radius $r$ and thickness $\delta r$ will bewhere $\sigma_{\rm DM}$ is dark matter velocity dispersion, $M(r)$ is total mass within radius $r$ and $v_c(r) = (GM/r)^{1/2}$ is the scale dependent virial velocity equivalent of the “circular velocity” for rotating disk.
Since the rotation curve has a roughly constant value vo at large radius, the ratio 2(T)/|(V)| will tend to the equilibrium value of 1 at large radius if opq*&vo.
Since the rotation curve has a roughly constant value $v_0$ at large radius, the ratio $2\langle T \rangle/|\langle V \rangle|$ will tend to the equilibrium value of $1$ at large radius if $\sigma_{\rm DM} \approx v_0$.
Using the minimumvalue for D from equation (3)), the maximum extent of the structureRmax will be D/(1—f.) and the average density will be where s=fmaxr/Tc.
Using the minimumvalue for $D$ from equation \ref{eq:roche}) ), the maximum extent of the structure$R_{max}$ will be $D/(1-f_r)$ and the average density will be where $s=r_{max}/r_c$.
Note, however, that this is a fractal structure with significant porosity.
Note, however, that this is a fractal structure with significant porosity.
So, the average density of any individualclump is higher than (p) by a factor of (1—nf?)(Rmax/Tmax)*.
So, the average density of any individualclump is higher than $\langle\rho\rangle$ by a factor of $(1-nf_r^3)(R_{max}/r_{max})^3$ .
Now, for global stability of the whole structure, Rmaz should be less than or equal to the radius within which the virial equilibrium is maintained.
Now, for global stability of the whole structure, $R_{max}$ should be less than or equal to the radius within which the virial equilibrium is maintained.
Average density within this virial radius Tvir Or T2990 Should be α΄200 times more than the critical density peri,=3H? /8xG.
Average density within this virial radius $r_{vir}$ or $r_{200}$ should be $\approx 200$ times more than the critical density $\rho_{crit} = 3H^2/8\pi G$ .
For a choice of model parameters n, f,, s and x, this will constrain thelower
For a choice of model parameters $n$ , $f_r$ , $s$ and $x$ , this will constrain thelower
interval. 3.2-10! 43107 s. indicated in Fig.
interval, $3.2\times 10^4$ $4.3\times 10^4$ s, indicated in Fig.
2 (hereafter this time interval is called “active” state and the rest ds “quiescent” state)
2 (hereafter this time interval is called “active” state and the rest is “quiescent” state).
We have verified that the position of the brightest X-ray source during the “active” period coincides with the position of source A. or the optical nucleus of NGC4395.
We have verified that the position of the brightest X-ray source during the “active” period coincides with the position of source A, or the optical nucleus of NGC4395.
This rules out the possibility that one of the nearby X-ray sources is responsible for the X-ray Daring.
This rules out the possibility that one of the nearby X-ray sources is responsible for the X-ray flaring.
Vhe normalized excess variance. Eepsq. was defined in Nandra et al (1997) as a measure of variability amplitude. (
The normalized excess variance, $\sigma^2_{\rm RMS}$, was defined in Nandra et al (1997) as a measure of variability amplitude. (
Note that Turner et al (1999) have pointed out an mistake in the formula quoted by Nandra et al (1997) for the error on ipe
Note that Turner et al (1999) have pointed out an mistake in the formula quoted by Nandra et al (1997) for the error on $\sigma^2_{\rm RMS}$.
We have used the corrected formula given by "Turner et al 1999.)
We have used the corrected formula given by Turner et al 1999.)
We computed ops for the 2:10 keV. flight curves and. found 0.203+0.066 for the SIS and 0.176=0.047 for the CIS.
We computed $\sigma^2_{\rm RMS}$ for the 2–10 keV light curves and found $0.203\pm 0.066$ for the SIS and $0.176\pm 0.047$ for the GIS.
Our lisht curve data are not sullicient to calculate a fully sampled. power spectrum by. standard. means.
Our light curve data are not sufficient to calculate a fully sampled power spectrum by standard means.
llowever we estimate the spectral properties. by two methods.
However we estimate the spectral properties by two methods.
First. we estimate the power spectrum using the algorithm: of Lomb (1976) which is specially designed. for irregularly sampled: data.
First, we estimate the power spectrum using the algorithm of Lomb (1976) which is specially designed for irregularly sampled data.
The result. shows the red-noise character tvpical of more luminous AGN.
The result shows the red-noise character typical of more luminous AGN.
Next. we make a quantitative comparison to other AGN. by calculating the normalized power spectral density (NPSDJ. following the methodology of Havashida et al (1998) at a small number of frequencies. which is presented in Fig.
Next, we make a quantitative comparison to other AGN, by calculating the normalized power spectral density (NPSD), following the methodology of Hayashida et al (1998) at a small number of frequencies, which is presented in Fig.
3.
3.
The error bars on each data point are too large to constrain the slope of the power spectrum.
The error bars on each data point are too large to constrain the slope of the power spectrum.