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Nevertheless. the results of this study can be useful in follow-up observations of exoplanet candidates discovered in transit surveys. by determining the parameters of especially the faint targets in a single observation based on the shape of its CCF bisector. | Nevertheless, the results of this study can be useful in follow-up observations of exoplanet candidates discovered in transit surveys, by determining the parameters of especially the faint targets in a single observation based on the shape of its CCF bisector. |
So far. most planet searches have been trying to smear or average out the effects of the stellar atmosphere (e.g.. ??).. | So far, most planet searches have been trying to smear or average out the effects of the stellar atmosphere \citep[e.g.,][]{dumusque+2011a,dumusque+2011b}. |
Only recently (2) has attempts been made to try to understand fully the effects of the stellar atmosphere with the aim of disentangling the effects. | Only recently \citep{boisse+2011} has attempts been made to try to understand fully the effects of the stellar atmosphere with the aim of disentangling the effects. |
This 1s of particular importance in the context of the coming planet finding instruments. where sub-m/s precision will call for ways to discern and ultimately disentangle the effects of the stellar atmosphere from the planetary signal (e.g..??).. | This is of particular importance in the context of the coming planet finding instruments, where sub-m/s precision will call for ways to discern and ultimately disentangle the effects of the stellar atmosphere from the planetary signal \citep[e.g., ESPRESSO for the VLT and later Codex for the E-ELT;][]{espresso,codex}.. |
10 keV to properly model and constrain the two spectral components. | 10 keV to properly model and constrain the two spectral components. |
The best-fitting model requires also the presence of a power-law (I« 2.9) at low energies and a eemission line at an energy of = 930 eV (EW= 50 eV). | The best-fitting model requires also the presence of a power-law $\Gamma\approx2.9$ ) at low energies and a emission line at an energy of $\approx$ 930 eV $EW\approx$ 50 eV). |
The measured 0.5-2 keV (2-10 Κεν) flux is zO23x107P (577x1077 s!)7... and the de-absorbed. rest-frame 2-10 keV luminosity is =8.410sv: we note. however. that this value is somehow uncertain. given the complex mmodeling of the data presented here. | The measured 0.5–2 keV (2–10 keV) flux is $\approx2.3\times10^{-13}$ $\approx5.7\times10^{-13}$ ), and the de-absorbed, rest-frame 2–10 keV luminosity is $\approx8.4\times10^{42}$; we note, however, that this value is somehow uncertain, given the complex modeling of the data presented here. |
An oobservation with a net exposure time of =65 ks has detected a faint ssource at this galaxy. | An observation with a net exposure time of $\approx 65$ ks has detected a faint source at this galaxy. |
The sspectrum (see Fig. | The spectrum (see Fig. |
A.2) shows a hard eexcess above 3 keV when a simple power-law is fitted. | A.2) shows a hard excess above 3 keV when a simple power-law is fitted. |
This hard ccomponent probably originates from. an. obscured active nucleus. | This hard component probably originates from an obscured active nucleus. |
When modeled by an absorbed power-law with photo index [=1.8. the absorbing column density Is estimated to be Nu=ghx1073 em. | When modeled by an absorbed power-law with photon index $\Gamma = 1.8$, the absorbing column density is estimated to be $N_{\rm H}=8^{+5}_{-4}\times 10^{23}$ $^{-2}$. |
The data have no sufficient quality te constrain a Fe K emission-line. | The data have no sufficient quality to constrain a Fe K emission-line. |
The soft X-ray component ca be described by a power-law of P=2.5. | The soft X-ray component can be described by a power-law of $\Gamma\approx 2.5$. |
While the origin of the soft eemission is unclear. assuming that no strong star formatio is taking place in this E/SO galaxy. it is likely extended | While the origin of the soft emission is unclear, assuming that no strong star formation is taking place in this E/S0 galaxy, it is likely extended |
dust erains for a large sample of AGNs. we focus on [Fe VIJAGOST aud [Ne VJA3126 again. | dust grains for a large sample of AGNs, we focus on [Fe $\lambda$ 6087 and [Ne $\lambda$ 3426 again. |
These two forbidden enission lines are thought to arise at similar regions since their critical deusities are similar (1.6&107 ? and 3.6.10° Ὁ, respectively). aud the ionization poteutials of Fe®! and Net! are also nearly the same (97.1 eV and 99.1 eV. respectively). | These two forbidden emission lines are thought to arise at similar regions since their critical densities are similar $1.6 \times 10^7$ $^{-3}$ and $3.6 \times 10^7$ $^{-3}$, respectively), and the ionization potentials of $^{6+}$ and $^{4+}$ are also nearly the same (97.1 eV and 99.1 eV, respectively). |
Iu addition to this advantage. it should be also au great advautage that these two enission ines are chough strong to be measured easily. | In addition to this advantage, it should be also an great advantage that these two emission lines are enough strong to be measured easily. |
Therefore. hey are useful to examine the dust abundances in the TINERs if we know some possible dependences of their enissivities on plivsical properties. even though they might © inappropriate to determine exact gas-phase elemental abuudances as noted by Ferguson ot al. ( | Therefore, they are useful to examine the dust abundances in the HINERs if we know some possible dependences of their emissivities on physical properties, even though they might be inappropriate to determine exact gas-phase elemental abundances as noted by Ferguson et al. ( |
1997). | 1997b). |
Iu this oper. we investigate the dependeuces of the two hieh-ionization enission lines on physical properties based on photoionization model calculations. | In this paper, we investigate the dependences of the two high-ionization emission lines on physical properties based on photoionization model calculations. |
We then examine how hese enisson lines can give constraints on the issue whether or not the TINERs are dusty. | We then examine how these emission lines can give constraints on the issue whether or not the HINERs are dusty. |
To study the properties of eas clouds in HINERs. we have compiled the data of cinissiou-line ffux ratios of [Fe vii|AGOST/|O HYA5007 aud [Ne v|A3126/]0 111|A5007 from the literature. | To study the properties of gas clouds in HINERs, we have compiled the data of emission-line flux ratios of [Fe $\lambda$ 6087/[O $\lambda$ 5007 and [Ne $\lambda$ 3426/[O $\lambda$ 5007 from the literature. |
The details of the data compilation are eiven by Nagao et al. ( | The details of the data compilation are given by Nagao et al. ( |
20015). | 2001b). |
The umuuber of compiled ohjects is 58: 31 tvpe 1. AGNs and 21 tvpe 2 ACNs. | The number of compiled objects is 58; 34 type 1 AGNs and 24 type 2 AGNs. |
Tere we refer to Vérrou-Cetty Veérron (2001) for the ACN type of cach object. | Here we refer to Vérron-Cetty Vérron (2001) for the AGN type of each object. |
Note that. in this paper. the objects classified as type 1.0. 1.2. and 1.5 ACN (iucludiug narrow-line Sevtert 1 galaxies] by Vérrou-Cetty Vórron (2001) are included in “type 1 ACGN. and the objects classified as type La. 1.9. and 2.0 ACN are included in “type 2 AGN”. | Note that, in this paper, the objects classified as type 1.0, 1.2, and 1.5 AGN (including narrow-line Seyfert 1 galaxies) by Vérron-Cetty Vérron (2001) are included in “type 1 AGN”, and the objects classified as type 1.8, 1.9, and 2.0 AGN are included in “type 2 AGN”. |
Since we do not impose anv selection criteria upon our sample. this sample is ucither an wuiform nor complete one in any seuse. | Since we do not impose any selection criteria upon our sample, this sample is neither an uniform nor complete one in any sense. |
However. this does not affect the following discussion significantly because we are uot interested in statistical properties of NLR eas clouds. | However, this does not affect the following discussion significantly because we are not interested in statistical properties of NLR gas clouds. |
Namely. we focus on whether distribution of emission-line flux ratios is consistent with that expected for dust-free cases or dusty cases. | Namely, we focus on whether distribution of emission-line flux ratios is consistent with that expected for dust-free cases or dusty cases. |
The compiled euission-line fix ratios are given in Table 1. | The compiled emission-line flux ratios are given in Table 1. |
We adopt anu average value if au enission-liue flux ratio for a certain object is given in more than one paper published previously. | We adopt an average value if an emission-line flux ratio for a certain object is given in more than one paper published previously. |
As for the data of the enüssion-liue flux ratios. we do not make any reddening correction since it is often difficult to measure the fluxes of narrow coniponeuts of Baluer lines for type LAGNs (sec. ce. Nagao et al. | As for the data of the emission-line flux ratios, we do not make any reddening correction since it is often difficult to measure the fluxes of narrow components of Balmer lines for type 1 AGNs (see, e.g., Nagao et al. |
200110). | 2001b). |
Effects of the dust extinction are discussed when uecessiry. | Effects of the dust extinction are discussed when necessary. |
Tn Figure 1. we show the frequency distributions of the two cliission-line fix ratios. [Ne v]A3126/]O ΠΠΙΑΡΟΕΤ aud [Fe vir|AG608T/[O. LJA5007. for both the type 1 AGNs aud the type 2 ACNs in our sample. | In Figure 1, we show the frequency distributions of the two emission-line flux ratios, [Ne $\lambda$ 3426/[O $\lambda$ 5007 and [Fe $\lambda$ 6087/[O $\lambda$ 5007, for both the type 1 AGNs and the type 2 AGNs in our sample. |
The average aud the 1-0 standard deviation of these two emission-liue fux ratios are 0.191 + O11L and 0.089 + 0.079 for the type 1 ACNs and 0.066 + 0.012 and 0.020 + 0.015 for the type 2 ACNs. respectively, | The average and the $\sigma$ standard deviation of these two emission-line flux ratios are 0.191 $\pm$ 0.114 and 0.089 $\pm$ 0.079 for the type 1 AGNs and 0.066 $\pm$ 0.042 and 0.020 $\pm$ 0.015 for the type 2 AGNs, respectively. |
The comparison of these results between the type Laud type 2 ACNs clearly suggests that the type 1l AGNs exhibit systematically larger ratios of both [Ne VJA3126//O 11]A5007 and [Fe vuJAGOS7/[O 11]A5007 than the type 2 Δέλλα | The comparison of these results between the type 1 and type 2 AGNs clearly suggests that the type 1 AGNs exhibit systematically larger ratios of both [Ne $\lambda$ 3426/[O $\lambda$ 5007 and [Fe $\lambda$ 6087/[O $\lambda$ 5007 than the type 2 AGNs. |
, We apply the Noluoeorov-Simirnov (1ν-S) statistical test on the data to estimate the siguificauce of the differences in the two cuissiou-line flux ratios between the type 1 ACGNs aud the type 2 AGNs. | We apply the Kolmogorov-Smirnov (K-S) statistical test on the data to estimate the significance of the differences in the two emission-line flux ratios between the type 1 AGNs and the type 2 AGNs. |
The | The |
The net accretion rate Aly onto the black hole need not equal the viscous intlow rate Mas abo LOO pe in equation (1)) if there are significant outllows on scales 100 pe (roughly our resolution): we return to this point in 2.4. | The net accretion rate $\dot M_{\rm in}$ onto the black hole need not equal the viscous inflow rate $\dot M_{\rm visc}$ at $\sim 100$ pc in equation \ref{eqn:Mdvisc}) ) if there are significant outflows on scales $\lesssim 100$ pc (roughly our resolution); we return to this point in 2.4. |
In the radiation pressure feedback model introduced in DOM. the accretion luminosity is coupled back into the surrounding eas by adding a total force which is shared equally by all SPLL particles inside 2.,.... | In the radiation pressure feedback model introduced in DQM, the accretion luminosity is coupled back into the surrounding gas by adding a total force which is shared equally by all SPH particles inside $R_{acc}$. |
Phe added force is directed racially away from the BIL particle. | The added force is directed radially away from the BH particle. |
llere 7 is à free parameter of the model. representing the optical depth to infrared. (LR) radiation. in the. nuclear region. | Here $\tau$ is a free parameter of the model, representing the optical depth to infrared (IR) radiation in the nuclear region. |
Our [iducial choice for 7 in this work is 20. | Our fiducial choice for $\tau$ in this work is 20. |
DON, showed that 7~20 was required to reproduce the observed normalization of the AlpyH o@ correlation in their simulations. | DQM showed that $\tau \sim 20$ was required to reproduce the observed normalization of the $M_{\rm BH}-\sigma$ correlation in their simulations. |
In addition to the radiation pressure feedback described by equation (3)). accretion onto the black hole can explicitly drive a wind at radii well beneath our resolution. | In addition to the radiation pressure feedback described by equation \ref{eqn:momdep}) ), accretion onto the black hole can explicitly drive a wind at radii well beneath our resolution. |
Our treatment of such a wind is motivated. in particular by observations of BAL quasars. which have ~10.000kms outIlows launched from near the broad line region at ~0.1 pe (eg. 2)). | Our treatment of such a wind is motivated in particular by observations of BAL quasars, which have $\sim 10,000 \kms$ outflows launched from near the broad line region at $\sim 0.1$ pc (e.g., \citealt{murray1995}) ). |
In our model. the ACGN winds at small raclii carry a momentum fux given by where zr, is a further parameter of the model representing the total momentum flux in the wind. | In our model, the AGN winds at small radii carry a momentum flux given by where $\tau_w$ is a further parameter of the model representing the total momentum flux in the wind. |
The wind is launched at a fixed speed. eo. | The wind is launched at a fixed speed $v_w$. |
Phus. the rate at which mass is added to the wind is given by where we have linked the luminosity of the black hole to the accretion rate as in equation. (2)). | Thus, the rate at which mass is added to the wind is given by where we have linked the luminosity of the black hole to the accretion rate as in equation \ref{eqn:LfromMdot}) ). |
In the presence of significant AGN winds. not all of the material entering the nuclear region. Alpine actually reaches the black hole. | In the presence of significant AGN winds, not all of the material entering the nuclear region, $\dot{M}_{visc}$, actually reaches the black hole. |
Instead. the net aceretion rate into the black hole. AT;,. must be reduced. by the mass of the outllow (sec 2)): Mi,ΔίονM which implies that the truc black hole accretion rate is given by Note that equation (6)). not equation (11). determines the AGN luminosity and thus the magnitude of the feedback in equations (3)) ancl (4)). | Instead, the net accretion rate into the black hole, $\dot{M}_{in}$, must be reduced by the mass of the outflow (see \citealt{ostriker10}) ): $\dot{M}_{in} = \dot{M}_{visc} -
\dot{M}_w$, which implies that the true black hole accretion rate is given by Note that equation \ref{eqn:Mdin}) ), not equation \ref{eqn:Mdvisc}) ), determines the AGN luminosity and thus the magnitude of the feedback in equations \ref{eqn:momdep}) ) and \ref{eqn:pdwind}) ). |
To deposit the wind. momentum. we give. kicks to particles inside {μοι | To deposit the wind momentum, we give kicks to particles inside $R_{acc}$. |
Particles receiving a kick are selecte stochastically with each particle having an equal probability of being added to the wind. | Particles receiving a kick are selected stochastically with each particle having an equal probability of being added to the wind. |
Phe probability of a kick is chosen to ensure that the time average of the mass kick in a given timestep Af is given hy AlyAt. | The probability of a kick is chosen to ensure that the time average of the mass kicked in a given timestep $\Delta t$ is given by $\dot M_w \Delta t$. |
Vhe kick is implemented by adding a velocity 0; to the current velocity of cach selected: particle. | The kick is implemented by adding a velocity $v_w$ to the current velocity of each selected particle. |
Phe direction of this imparte velocity is chosen to be more heavily weighted toward the surrounding clisk: if @ is the angle between the impartec momenttun ancl the disk normal. the probability clistribution over ar=cos@ is given by pir)=31)/4. | The direction of this imparted velocity is chosen to be more heavily weighted toward the surrounding disk: if $\theta$ is the angle between the imparted momentum and the disk normal, the probability distribution over $x = \cos{\theta}$ is given by $p(x) = 3(1 - x^2)/4$. |
For this purpose. the disk normal is defined to be the directionof the total orbital angular momentum about the black hole of all the eas particles inside 2... | For this purpose, the disk normal is defined to be the directionof the total orbital angular momentum about the black hole of all the gas particles inside $R_{acc}$. |
This modest equatorial bias in the wind direction is motivated by models of BAL quasars (2).. | This modest equatorial bias in the wind direction is motivated by models of BAL quasars \citep{murray1995}. |
Nonetheless. this choice is not that critical: the ellicaev of the feedback is largely determined by the outllow momentum/energv [lux that is directed. within the solid angle subtended by the surrounding ISAL | Nonetheless, this choice is not that critical: the efficacy of the feedback is largely determined by the outflow momentum/energy flux that is directed within the solid angle subtended by the surrounding ISM. |
For our fiducial model 40% of the momentum Lux is directed within one scale-height of the disk at ~LOO pc. | For our fiducial model, $\sim 40 \%$ of the momentum flux is directed within one scale-height of the disk at $\sim 100$ pc. |
Models with isotropic kicks require slightly large values of 7, to give results similar to our fiducial model because less of the feedback is directed towards the surrounding ISM. | Models with isotropic kicks require slightly large values of $\tau_w$ to give results similar to our fiducial model because less of the feedback is directed towards the surrounding ISM. |
Our implementation of wind. feedback. adds two aclelitional input parameters to the simulation. | Our implementation of wind feedback adds two additional input parameters to the simulation. |
The first. 64. cleseribes the launch speed of the winds while the second. 7. describes the total momentum flux in the wind. | The first, $v_w$, describes the launch speed of the winds while the second, $\tau_w$ , describes the total momentum flux in the wind. |
Our cefault wind speed is motivated in part by observations of BAL quasars and theoretical models of the origin of such winds via line driving in the accretion disk at ~0.1 pe (though our model should not be taken as a literal implementation of this physics). | Our default wind speed is motivated in part by observations of BAL quasars and theoretical models of the origin of such winds via line driving in the accretion disk at $\sim 0.1$ pc (though our model should not be taken as a literal implementation of this physics). |
Observed wind velocities are 10.000kms | Observed wind velocities are $\sim 10,000 \kms$. |
AMoclels of line criving lead to momentum Iluxes in the wind ol Líc.le. Tal because the lines do not. typically completely cover the continuum (2)... | Models of line driving lead to momentum fluxes in the wind of $\lesssim L/c$, i.e., $\tau_w \lesssim 1$ because the lines do not typically completely cover the continuum \citep{murray1995}. |
As we show below. however. larger values of τι ave τουτος for GN outllows to have a significant. effect on the gas dynamics and. star formation history in galactic nuclei. | As we show below, however, larger values of $\tau_w$ are required for AGN outflows to have a significant effect on the gas dynamics and star formation history in galactic nuclei. |
Observationallv. momentum Iuxes are cillicult to infer in mostcases because of ambiguities in. the absolute densitv/radial scale at which the absorption occurs. | Observationally, momentum fluxes are difficult to infer in mostcases because of ambiguities in the absolute density/radial scale at which the absorption occurs. |
In a handful of low-ionization BALs (in particular. FeLoBALs). this degeneracy has been broken. implying momentum Iuxes 90.391ο. οι te00.5.5 (with significant uncertainties: see ?2??7)). | In a handful of low-ionization BALs (in particular, FeLoBALs), this degeneracy has been broken, implying momentum fluxes $\sim
0.3-5 L/c$, i.e., $\tau_w \sim 0.3-5$ (with significant uncertainties; see \citealt{moe09, bautista10,dunn10,claude11}) ). |
Ht remains uncertain whether these values of 7, are representative of the entire BAL-quasar population. | It remains uncertain whether these values of $\tau_w$ are representative of the entire BAL-quasar population. |
We ake T,&5 as our fiducial value but also explore the range ]Xore< 10. | We take $\tau_w \simeq 5$ as our fiducial value but also explore the range $1\le \tau_w \le 10$ . |
Our ducial value of τι=5 was chosen for hree reasons that will become clear later in the paper: (1) he ACGN wind then leads to a ealaxy-scale outflow. that significantly influences the surrounding ISM dynamics. (2) he final DII mass in the simulations is reasonably consistent withthe observed AJpyH0 relation for τι2ο 10. ( | Our fiducial value of $\tau_w = 5$ was chosen for three reasons that will become clear later in the paper: (1) the AGN wind then leads to a galaxy-scale outflow that significantly influences the surrounding ISM dynamics, (2) the final BH mass in the simulations is reasonably consistent withthe observed $M_{\rm BH}-\sigma$ relation for $\tau_w \simeq 5-10$ . ( |
3) Observations of high speedatomic and molecular outllows in ocal ULIRGs suggest 7,10 (7777).. | 3) Observations of high speedatomic and molecular outflows in local ULIRGs suggest $\tau_w \sim 10$ \citep{feruglio10,chung11,rupke11,sturm11}. . |
We discuss physical | We discuss physical |
well studied SNR, for which a relatively accurate determination of parameters such as distance, shock velocity and age, is available. | well studied SNR, for which a relatively accurate determination of parameters such as distance, shock velocity and age, is available. |
SN 1006 is an historical SNR with age ~1000 yr and believed to be in the Sedov phase of its evolution. | SN 1006 is an historical SNR with age $\approx 1000$ yr and believed to be in the Sedov phase of its evolution. |
After reviewing the preexistingliterature, ? concluded that the distance to SN 1006 lies in the range 1.4—2 kpc. | After reviewing the preexistingliterature, \cite{Dubner2002} concluded that the distance to SN 1006 lies in the range 1.4–2 kpc. |
More recently, a slightly larger distance of 2.18 kpc was inferred by ?,, who compared the measured optical proper motion with the expansion velocity of 2890+100 km/s determined from Ha lines observations. | More recently, a slightly larger distance of $2.18$ kpc was inferred by \cite{Winkler2003}, who compared the measured optical proper motion with the expansion velocity of $2890 \pm 100$ km/s determined from $\alpha$ lines observations. |
This velocity is in a good agreement with analogous velocity measurements by ? which give v—2200—3500 km/s. At this distance the SNR apparent size 15’ (?)) translates into a shock radius equal to rsnock=9.6 pc. | This velocity is in a good agreement with analogous velocity measurements by \cite{Smith1991} which give $v=2200-3500$ km/s. At this distance the SNR apparent size 15' \citealt{Rothenflug2004}) ) translates into a shock radius equal to $r_{shock}=9.6$ pc. |
A value for the ambient gas density equal to n=0.15—0.25 cm-? (in the north-western region) has been obtained by ? from XMM X-ray observations. | A value for the ambient gas density equal to $n=0.15-0.25$ $^{-3}$ (in the north-western region) has been obtained by \cite{Acero2007} from XMM X-ray observations. |
A slightly higher values for the density n=0.25—0.4 cm? was obtained from Ha lines observations with HST (?)). | A slightly higher values for the density $n=0.25-0.4$ $^{-3}$ was obtained from $\alpha$ lines observations with HST \citealt{Raymond2007}) ). |
By assuming that the evolution of SN 1006 can be described by the Chevalier self-similar model described in Sec. D], | By assuming that the evolution of SN 1006 can be described by the Chevalier self–similar model described in Sec. \ref{sec:SNR dynamics}, |
it is possible to derive a supernova explosion energy equal to Egy&2x10°! erg. | it is possible to derive a supernova explosion energy equal to $E_{SN} \approx 2 \times 10^{51}$ erg. |
Results from our model for SN 1006 are shown in Fig. [6], | Results from our model for SN 1006 are shown in Fig. \ref{fig:te-tp-sn1006}, |
where the electron and proton temperatures behind the SNR shock wave are plotted as afunction of the remnant age. | where the electron and proton temperatures behind the SNR shock wave are plotted as afunction of the remnant age. |
For the actual age of the remnant, t~1000 yr, ambient gas density n=0.2 οπιὉ, explosion energy E=2-10?! erg and w=0.1, our model gives a value of the proton temperature equal to Ty~17 keV, and an electron temperature equal to Τε=0.6 keV. The observational values are from ? for the proton and electron temperatures are also indicated in Fig. [6], | For the actual age of the remnant, $t \approx 1000$ yr, ambient gas density $n=0.2$ $^{-3}$ , explosion energy $E=2\cdot 10^{51}$ erg and $w=0.1$, our model gives a value of the proton temperature equal to $T_p \sim 17$ keV, and an electron temperature equal to $T_e = 0.6$ keV. The observational values are from \cite{Vink2003} for the proton and electron temperatures are also indicated in Fig. \ref{fig:te-tp-sn1006}, , |
together with the error bars (3c for the proton temperature and confidence level for the electron temperature). | together with the error bars $3 \sigma$ for the proton temperature and confidence level for the electron temperature). |
A bit lower proton temperature T,=1.8:105K16keV was obtained by ?.. | A bit lower proton temperature $T_p=1.8\cdot 10^8K\approx 16keV$ was obtained by \cite{Korreck2004}. |
It is clear from Fig. | It is clear from Fig. |
|ó| that our model predicts a proton temperature which is in agreement with observations, while the prediction for the electron temperature falls a factor of 2-3 below the observational value of Τεz1.5+0.2 keV (?),, (MM). | \ref{fig:te-tp-sn1006} that our model predicts a proton temperature which is in agreement with observations, while the prediction for the electron temperature falls a factor of 2–3 below the observational value of $T_e\approx 1.5\pm 0.2$ keV \citep{Vink2003}, (XMM). |
A similar value for the electron temperature of T,=1.5—1.7 keV has been independently measured by ? with XMM. | A similar value for the electron temperature of $T_e = 1.5 - 1.7$ keV has been independently measured by \cite{Acero2007} with XMM. |
Nevertheless there is some doubt in the electron temperature measurements. | Nevertheless there is some doubt in the electron temperature measurements. |
Lower (0.6-0.7keV) values were obtained by ? with CHANDRA for NW-1 region, while other NW regions appear to be ejecta dominated. | Lower (0.6-0.7keV) values were obtained by \cite{Long2003} with CHANDRA for NW-1 region, while other NW regions appear to be ejecta dominated. |
Since XMM has worse spatial resolution than CHANDRA, ejecta dominated regions can contribute significantly to its spectra. | Since XMM has worse spatial resolution than CHANDRA, ejecta dominated regions can contribute significantly to its spectra. |
Moreover, in ? itis shown that including synchrotron emission component can reduce observed electron temperature. | Moreover, in \cite{Acero2007} it is shown that including synchrotron emission component can reduce observed electron temperature. |
If density and velocity vary in the range n—0.15 cm-?, v=2400—3500 km/s, respectively, we estimated the corresponding range of electron temperatures predictions from the model to be Τε=0.5—0.9 keV. This values are in the agreement with ones from the model in ?,, see their Fig. | If density and velocity vary in the range $n=0.15-0.5$ $^{-3}$, $v=2400-3500$ km/s, respectively, we estimated the corresponding range of electron temperatures predictions from the model to be $T_e=0.5-0.9$ keV. This values are in the agreement with ones from the model in \cite{Vink2003}, see their Fig. |
2. | 2. |
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