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Du other words. If 9; expresses the nuniber of particles of species i per ores per cn. then a;=στπο]. | In other words, if $\hat{n}_i$ expresses the number of particles of species $i$ per ergs per $^3$, then $n_i = \hat{n}_i(\sth R)(\melec c^2)$. |
Tie is rorlualized with respect to the photon crossine/escapeine from the source. £4,= Re. | Time is normalized with respect to the photon crossing/escapetime from the source, $t_{\rm cr}=R/c$ . |
Thus. 7. which appears iu the sinetic equatious. is dimensionless aud equals T= n2 The operators Q and £ denote injection aud osses. respectively, | Thus, $\tau$ , which appears in the kinetic equations, is dimensionless and equals $\tau=\frac{ct}{R}$ The operators $\cal{Q}$ and $\cal{L}$ denote injection and losses, respectively. |
The only processes that we takeiuto account are 55 annihilation and svuchrotron cooling of the produced pais | The only processes that we takeinto account are $\gamma \gamma$ annihilation and synchrotron cooling of the produced pairs. |
The svuchrotron οκνττν is | The synchrotron emissivity is |
ἡ0140. | $\eta \sim 140$. |
The total flix is shown by the thin solid line. | The total flux is shown by the thin solid line. |
In this case. we need to assume larger ISAL turbulence. | In this case, we need to assume larger ISM turbulence. |
In this section. we assume (he values of parameters wilh which we have shown the example cases in (le previous section. and estimate the X-ray. ancl radio afterelows. | In this section, we assume the values of parameters with which we have shown the example cases in the previous section, and estimate the X-ray and radio afterglows. |
The exiinclion correction a, is unity lor X-ray and radio afterglows. and hence Finer1.6 mJv (radio ancl X-ray). | The extinction correction $a_\nu$ is unity for X-ray and radio afterglows, and hence $F_{\nu,max,f} \sim 1.6$ mJy (radio and X-ray). |
X-ray Afterglow: Since the N-rav. band. ~5keV is well above the tvpical frequency of the reverse shock emission. the contribution from the reverse shock to the N-rav. band is negligible. | X-ray Afterglow: Since the X-ray band $\sim 5$ keV is well above the typical frequency of the reverse shock emission, the contribution from the reverse shock to the X-ray band is negligible. |
The X-ray afterglow should be described only by the forward shock emission. | The X-ray afterglow should be described only by the forward shock emission. |
The huninosity in X-ray band should decrease as [2PE4tip=2,4) or) FP(p=2.2). | The luminosity in X-ray band should decrease as $t^{(2-3p)/4} \sim t^{-1.3}(p=2.4)$ or $t^{-1.15}(p=2.2)$. |
The Chandra X-ray. observatory observed (he alterglow for a total exposure of STksec. beginning al Oct 5 8:55 UT (Sako and Harrison 2002). | The Chandra X-ray observatory observed the afterglow for a total exposure of 87ksec, beginning at Oct 5 8:55 UT (Sako and Harrison 2002). |
The count rate decrease with a power law slope of —1.02:0.2. | The count rate decrease with a power law slope of $-1.0 \pm
0.2$. |
The mean 2-10 keV. X-ray flux is ~4.3x10.P eres 7s 1. | The mean 2-10 keV X-ray flux is $\sim 4.3 \times10^{-13}$ ergs $^{-2}$ $^{-1}$. |
We estimate the 5keVT flux at the observational. mean time. 1.36⋅⋅ days ~3.1.x1015 eres »7? ! forB p24 and 6.4x10.P eres ? ! for p=2.2. | We estimate the 5keV flux at the observational mean time 1.36 days $\sim 3.1\times10^{-13}$ ergs $^{-2}$ $^{-1}$ for $p=2.4$ and $\sim 6.4\times10^{-13}$ ergs $^{-2}$ $^{-1}$ for $p=2.2$. |
Our estimates are in a good agreement with the observations. | Our estimates are in a good agreement with the observations. |
Radio Afterglow: The forward shock emission in radio band ~10 Gllz increases as [7 until the flux reaches to the maximum ~1.6 mJv al ~80 days for p=2.4 (dashed dotted line in fig?) or at ~50 davs lor p=2.2. | Radio Afterglow: The forward shock emission in radio band $\sim 10$ GHz increases as $t^{1/2}$ until the flux reaches to the maximum $\sim 1.6$ mJy at $\sim 80$ days for $p=2.4$ (dashed dotted line in fig2) or at $\sim 50$ days for $p=2.2$. |
After the typical [requeney 14, crosses (he radio band. the reverse shock emission decays as ~/7 (dashed line for p= 2.4). | After the typical frequency $\nu_{m,r}$ crosses the radio band, the reverse shock emission decays as $\sim t^{-2}$ (dashed line for $p=2.4$ ). |
M low frequencies and early limes. self absorption takes an important role ancl significantly reduces ihe flux. | At low frequencies and early times, self absorption takes an important role and significantly reduces the flux. |
A simple estimate of the maximal flux (dotted line for p= 2.4) is the emission from the black body with the reverse shock temperature (IXobavashi Sari 2000). | A simple estimate of the maximal flux (dotted line for $p=2.4$ ) is the emission from the black body with the reverse shock temperature (Kobayashi Sari 2000). |
The thick and (hin solid line depicts the total {lus for p—2.4 and lor p=2.2. respectively. | The thick and thin solid line depicts the total flux for $p=2.4$ and for $p=2.2$, respectively. |
Since the observations (circles) are done in various frequencies. we scaled the observed. value to the expected value at. 106 by using a spectral slope of 1.1. | Since the observations (circles) are done in various frequencies, we scaled the observed value to the expected value at $10$ GHz by using a spectral slope of $1$ . |
This burst might also cause a bright radio flare 1 mJv around ~0.5 day as observed in GRB 990123. | This burst might also cause a bright radio flare $\sim 1$ mJy around $\sim 0.5$ day as observed in GRB 990123. |
When we fit the | When we fit the |
We have shown that the thiourea functional eroup. associated with various carbonaceous structures. has one or two strong emission lines in a spectral range of ~ Lynn. within the 21-52 band emitted by a umuber of pre-planetary uebulae. | We have shown that the thiourea functional group, associated with various carbonaceous structures, has one or two strong emission lines in a spectral range of $\sim$ 4 $\mu$ m, within the $\mu$ m band emitted by a number of pre-planetary nebulae. |
The combination of nitrogen aud sulphur in thiourea is the esseutial source of Cluission iu this model: the baud disappears if these species are replaced by carbon. | The combination of nitrogen and sulphur in thiourea is the essential source of emission in this model: the band disappears if these species are replaced by carbon. |
These two clements are part of the ubiquitous CITONS family because of their high chemical activity. | These two elements are part of the ubiquitous CHONS family because of their high chemical activity. |
Thiourea may therefore readily forma in space. aud be found as an independent molecule or as a peripheral eroup attached to carbouaceous structures believed to be abundant ii space. | Thiourea may therefore readily form in space, and be found as an independent molecule or as a peripheral group attached to carbonaceous structures believed to be abundant in space. |
Iu all cases. it carries a strong IR Lue near the molecular thiourea line. which is the strougest and thus determines the peak of the baud. | In all cases, it carries a strong IR line near the molecular thiourea line, which is the strongest and thus determines the peak of the band. |
Obvioush. no single structure can exhibit the required spectrum. for cach ouly contributes discrete lines which camunot be broadened enough by usual broadening mcchanisins. | Obviously, no single structure can exhibit the required spectrum, for each only contributes discrete lines which cannot be broadened enough by usual broadening mechanisms. |
Twelve structures have beeu selected here. but their list is far from begC» exhaustive: thev are only iuteuded as exaniples of the generic thiourea class. | Twelve structures have been selected here, but their list is far from being exhaustive; they are only intended as examples of the generic thiourea class. |
The ehienical software used hiere also allows to determine the types of modal vibrations which cary the lines of interest: this helps designing new structures to fill the wide bands observed im the sky. | The chemical software used here also allows to determine the types of modal vibrations which cary the lines of interest; this helps designing new structures to fill the wide bands observed in the sky. |
- Using interpolation and smoothing between the coucatenated discrete. lines of the selected structures. we produced svuthetic spectra which exhibit a prominent. asvuumnetrie. feature between I8 aud 25 gan. with ανακια points at 19.6 aud 21.9 gan. very near the observed values. | Using interpolation and smoothing between the concatenated discrete lines of the selected structures, we produced synthetic spectra which exhibit a prominent, asymmetric, feature between 18 and 25 $\mu$ m, with half-maximum points at 19.6 and 21.9 $\mu$ m, very near the observed values. |
However. the peak is 0.6 pan vedward of the observed average. | However, the peak is 0.6 $\mu$ m redward of the observed average. |
The astronomical 21-41. feature extends rechward to merge with the other. prominent 30-/24: band. | The astronomical $\mu$ m feature extends redward to merge with the other, prominent $\mu$ m band. |
It is found that the main characters of this band cau be modelled by the combined spectra of: a) aliphatic chains. made of CIT; eroups. oxvgen bridges and ΟΠ eroups. which provide the 30-4221: ciission: b) siiall. mostly lear. aromatic structures. which contribute to raise the red wing of the 21-422. baud and fill the space between the two main features. | It is found that the main characters of this band can be modelled by the combined spectra of: a) aliphatic chains, made of $_{2}$ groups, oxygen bridges and OH groups, which provide the $\mu$ m emission; b) small, mostly linear, aromatic structures, which contribute to raise the red wing of the $\mu$ m band and fill the space between the two main features. |
The concatenated spectral lines of ten of these structures form a stroug band between 23 and 38 qun. The omission of oxvseu in such structures all but extinguishes the 30-4 enusson. | The concatenated spectral lines of ten of these structures form a strong band between 23 and 38 $\mu$ m. The omission of oxygen in such structures all but extinguishes the $\mu$ m emission. |
The fact that these carriers do not involve. aud are likely more abundant than. thiourea derivatives eusures that the 30-420. feature can still be preseut in the absence of the 21-4211 feature. as observed. | The fact that these carriers do not involve, and are likely more abundant than, thiourea derivatives ensures that the $\mu$ m feature can still be present in the absence of the $\mu$ m feature, as observed. |
Combining the discrete lines of the 22 selected structures in different proportions. interpolating aud smoothing. we produced 2 svuthetic spectra which purport | Combining the discrete lines of the 22 selected structures in different proportions, interpolating and smoothing, we produced 2 synthetic spectra which purport |
(fy.bs) with E,=50 keV and E»=300 keV. For an object al redshift z. the observed energy range (f,.£5) originates in the range (I4(12-z).Es(124-:)) in the objects rest frame. whereas the luminosity refers to the range (£4.F5). | $(E_1,E_2)$ with $E_1 = 50$ keV and $E_2 = 300$ keV. For an object at redshift $z$, the observed energy range $(E_1,E_2)$ originates in the range $(E_1(1+z),E_2(1+z))$ in the object's rest frame, whereas the luminosity refers to the range $(E_1,E_2)$. |
The IX-term is the ratio of the rest frame energies raciated in the (wo ranges. The Dand photon spectrum is ususally described in terms of a break energy. £j. | The K-term is the ratio of the rest frame energies radiated in the two ranges, The Band photon spectrum is ususally described in terms of a break energy $E_0$ . |
Here we use a Band spectrum NCE.Ey...) where Ey(sp)=(2+a)Ly(sp) assuming that 2«—2 (Banelοἱal.1993). | Here we use a Band spectrum $N(E,E_{pk},\alpha,\beta)$ where $E_{pk}(sp)=(2 + \alpha)E_0(sp)$ assuming that $\beta < -2$ \citep{ban93}. |
. We adopt constant. values of a=—0.8 and 9=—2.6 for reasons that will be discussed in Section 5. | We adopt constant values of $\alpha = -0.8$ and $\beta = -2.6$ for reasons that will be discussed in Section 5. |
The peak flux P(L.z) observed for a GRB of Iuminositv.L at redshift 2 is where (c/H3).1(z) is the bolometric Iuminosity distance. | The peak flux $P(L,z)$ observed for a GRB of luminosity$L$ at redshift $z$ is where $(c/H_0)A(z)$ is the bolometric luminosity distance. |
We use (he cosmological parameters Hyτο kms ! !. Q4,—0.3. and Q4—0.7. | We use the cosmological parameters $H_0 = 70~$ km $^{-1}$ $^{-1}$ , $\Omega_M = 0.3$, and $\Omega_{\Lambda} = 0.7$. |
The integral peak flux distribution for GhRDs of spectral class sp is. where z(L.P.sp) is derived. [romequation (4). V(z) is the comoving volume and the term (1+2)! represents the time dilation. | The integral peak flux distribution for GRBs of spectral class $sp$ is, where $z(L,P,sp)$ is derived fromequation (4), $V(z)$ is the comoving volume and the term $(1+z)^{-1}$ represents the time dilation. |
With this formulation. it is straightforward to derive the differential source counts dN(>P.sp)/dP. as well as the average values of/Vinas-Poss . 03. elc. | With this formulation, it is straightforward to derive the differential source counts $dN(>P,sp)/dP$, as well as the average values of, $\alpha_{23}$ , etc. |
For a given rate funelion the procedure to iterate the huninoxity function is as follows. | For a given rate function the procedure to iterate the luminoxity function is as follows. |
Assume starting values for the central huninosity. Z.(sp) and the rest frame Band peak energy E,(sp). | Assume starting values for the central luminosity $L_c(sp)$ and the rest frame Band peak energy $E_{pk}(sp)$. |
The differential source counts together with —(D/Pj,,)*? produce the expected values of »Gp)forlheGUS BADcalalog. | The differential source counts together with $= (P/P_{lim})^{-3/2}$ produce the expected values of $(sp)$ for the GUSBAD catalog. |
Similarly. | Similarly, the expected values of $(sp)$ are obtained by weighting $/(1+z)$ with the differential source counts. |
theexpe lt i | The iteration isrepeated until the expected values of $(sp)$ and $(sp)$ match the observed ones given in Table 2. |
s worth noting that given aad (he shape of the five spectral luminosity ΠαπΙου. the procedure leads lor each spto a single value lor L.(sp) and Γρ). | It is worth noting that given and the shape of the five spectral luminosity functions, the procedure leads for each $sp$to a single value for $L_c(sp)$ and $E_{pk}(sp)$ . |
The primary unknown is the densityfiction ). | The primary unknown is the densityfunction . |
. We will use cldata to test various forms of R(z).. see Sec. | We will use data to test various forms of , see Sec. |
!5. | 5. |
aggregates are very fluffy, open structures. | aggregates are very fluffy, open structures. |
If one would draw a circumscribing sphere around the aggregate the vacuum fraction inside this sphere would be very large. | If one would draw a circumscribing sphere around the aggregate the vacuum fraction inside this sphere would be very large. |
in order to do computations for different aggregate sizes one might simply increase the size of the circumscribing sphere and perform Mie computations using the effective refractive index as given by Eq. (I0]). | in order to do computations for different aggregate sizes one might simply increase the size of the circumscribing sphere and perform Mie computations using the effective refractive index as given by Eq. \ref{eq:meffective}) ). |
We will refer to this method as the Aggregate Polarizability Mixing Rule (APMR). | We will refer to this method as the Aggregate Polarizability Mixing Rule (APMR). |
Now we have to determine how to get the required [πι and radius of the homogeneous sphere. | Now we have to determine how to get the required $f_\mathrm{fill}$ and radius of the homogeneous sphere. |
As is shown by (2006a) the most natural choice, namely using the circumscribing sphere, leads to an overestimate of the result of the fluffyness of the aggregate. | As is shown by \citet{2006A&A...445.1005M} the most natural choice, namely using the circumscribing sphere, leads to an overestimate of the result of the fluffyness of the aggregate. |
This is because the aggregate constituents are not randomly distributed in the volume circumscribed by this sphere. | This is because the aggregate constituents are not randomly distributed in the volume circumscribed by this sphere. |
Therefore, we have to choose a somewhat smaller radius which also somehow takes into account the structure of the aggregates. | Therefore, we have to choose a somewhat smaller radius which also somehow takes into account the structure of the aggregates. |
A choice, which leads to excellent results as shown below, is to use the radius of gyration of the aggregate. | A choice, which leads to excellent results as shown below, is to use the radius of gyration of the aggregate. |
In this equation r; is the location of constituent i, and ro is the center of mass of the aggregate. | In this equation $\vec{r}_i$ is the location of constituent $i$, and $\vec{r}_0$ is the center of mass of the aggregate. |
For fractal aggregates with fractal dimension D; the radius of gyration can be expressed as (Filippovetal2000) where y is a constant depending only on the size of the constituents and on the so-called fractal prefactor. | For fractal aggregates with fractal dimension $D_f$ the radius of gyration can be expressed as \citep{Filippov}
where $\gamma$ is a constant depending only on the size of the constituents and on the so-called fractal prefactor. |
By comparing the gyration radii computed using Eqs. ᾖᾖΤ)) | By comparing the gyration radii computed using Eqs. \ref{eq:rg}) ) |
and (12)) one finds that the aggregates we use are best represented using Dy=2.82 and y= 2.44. | and \ref{eq:rg2})) one finds that the aggregates we use are best represented using $D_f=2.82$ and $\gamma=2.44$ . |
The filling factor is now defined by | The filling factor is now defined by |
of the planet is given by qo. | of the planet is given by $\varphi_0$. |
The value of yo is determined by the relative velocity of the planetary and stellar coronal material. | The value of $\varphi_0$ is determined by the relative velocity of the planetary and stellar coronal material. |
Figure | of ? illustrates the scenarios leading to the various shock orientations. | Figure 1 of \cite{Vidotto:2010p809} illustrates the scenarios leading to the various shock orientations. |
There are two limiting cases: an “ahead-shock” (yy O0) forms when the planet is embedded in the stellar corona and a "dayside-shock" (gj.- 90°) forms when the radial wind velocity is very much greater than the relative azimuthal velocity of the planet. | There are two limiting cases: an “ahead-shock" $\varphi_0\rightarrow0$ ) forms when the planet is embedded in the stellar corona and a “dayside-shock" $\varphi_0\rightarrow90^\circ$ ) forms when the radial wind velocity is very much greater than the relative azimuthal velocity of the planet. |
Here we use models for the stellar corona and wind (22). to obtain a value for the plasma density at the planet. | Here we use models for the stellar corona and wind \citep{Vidotto:2010p809,Vidotto:2011p803} to obtain a value for the plasma density at the planet. |
These models assume a typical solar base density of np~100m? (2) and either an isothermal hydrostatic corona or an isothermal thermally driven wind. | These models assume a typical solar base density of $n_0\sim10^8\textrm{cm}^{-3}$ \citep{Withbroe:1988p829} and either an isothermal hydrostatic corona or an isothermal thermally driven wind. |
We assume an adiabatic shock with a maximum compression ratio of 4. | We assume an adiabatic shock with a maximum compression ratio of 4. |
For an isothermal corona of temperature 7. the density of stellar material. Πρι. at a planet of orbital radius Aj, is given by Equation 8 of ?.. For the stellar wind case we use mass conservation and the momentum equation. to obtain values for the radial velocity. i and density. 7,4. | For an isothermal corona of temperature $T$, the density of stellar material, $n_{\rm obs}$, at a planet of orbital radius $R_{\rm orb}$ is given by Equation 8 of \cite{Vidotto:2010p809}, For the stellar wind case we use mass conservation $(nu_rr^2=\textrm{const})$ and the momentum equation, to obtain values for the radial velocity, $u_r$ and density, $n_{\rm obs}$. |
The plasma density can then be converted into a density of fully ionized magnesium using the relation. where. ny.7j is the ratio of Magnesium number density to Hydrogen number density which is derived from the metallicity of the host star (2).. | The plasma density can then be converted into a density of fully ionized magnesium using the relation, where, $n_{\rm Mg}/n_{\rm H}$ is the ratio of Magnesium number density to Hydrogen number density which is derived from the metallicity of the host star \citep{Hebb:2009p806}. |
For WASP-12. ma./ni=6.76x10? (2). | For WASP-12, $n_{\rm Mg}/n_{\rm H}=6.76\times10^{-5}$ \citep{Vidotto:2010p809}. |
From this density we can then find bow shock geometries and orientations that fit the observations of ?.. | From this density we can then find bow shock geometries and orientations that fit the observations of \cite{Fossati:2010p838}. |
To investigate whether the model presented by ?. is able to reproduce the data from the near-UV observations. we use Monte Carlo radiative transfer calculations to produce simulated light curves. | To investigate whether the model presented by \cite{Vidotto:2010p809} is able to reproduce the data from the near-UV observations, we use Monte Carlo radiative transfer calculations to produce simulated light curves. |
The parameters we adopt to match the WASP-12 system are: M,=M;. Rk,=1.79R, (where My and Ry are the mass and radius ofL.431 Jupiter). M,=1.35M. and A,=ΤΝ... | The parameters we adopt to match the WASP-12 system are: $M_p = 1.41 M_J$, $R_p = 1.79 R_J$ (where $M_J$ and $R_J$ are the mass and radius of Jupiter), $M_\star = 1.35 M_\odot$ and $R_\star = 1.57 R_\odot$. |
The host star is a late F type and the planet orbits in the equatorial plane with an impact parameter b=0.36R, (2). | The host star is a late F type and the planet orbits in the equatorial plane with an impact parameter $b= 0.36\,R_\star$ \citep{Hebb:2009p806}. |
The shocked material is considered to be at a distance ry, from the planet. with a thickness Ary, and an angular extent 2Ay. | The shocked material is considered to be at a distance $r_M$ from the planet, with a thickness $\Delta r_M$ and an angular extent $2\Delta \varphi$. |
The projected lateral extent of the shock is dependent on 73.qo and Ay. | The projected lateral extent of the shock is dependent on $r_M,\varphi_0$ and $\Delta\varphi$. |
The maximum distance between the planet and the projected lateral extent of the shock. X. can take the following forms Our simulated transit light curves are produced using a 3D Monte Carlo radiation transfer code (2).. | The maximum distance between the planet and the projected lateral extent of the shock, $X_M$ , can take the following forms Our simulated transit light curves are produced using a 3D Monte Carlo radiation transfer code \citep{Wood:1999p822}. |
The circumplanetary density structure is prescribed on a 3D spherical polar grid (coordinates r. 8. ) and is externally irradiated with Monte Carlo photon packets with a distribution that reproduces the spatial intensity distribution of a limb-darkened star. | The circumplanetary density structure is prescribed on a 3D spherical polar grid (coordinates $r$, $\theta$, $\varphi$ ) and is externally irradiated with Monte Carlo photon packets with a distribution that reproduces the spatial intensity distribution of a limb-darkened star. |
We assume a spherical planet and a darkening law such that the intensity / is given by where p=cos@2(1-7)7.0rLH is the radial distance into the stellar disk normalized to the stellar radius and /(0) is the emergent intensity at the centre of the star (2).. | We assume a spherical planet and a limb-darkening law such that the intensity $I$ is given by where $\mu = \cos\theta= (1-r^2)^{1/2}, 0\leq r\leq 1$ is the radial distance into the stellar disk normalized to the stellar radius and $I(0)$ is the emergent intensity at the centre of the star \citep{Mandel:2002p814}. |
The coethcients a, are chosen from ? to match the ii-band limb-darkening of the host star. | The coefficients $a_n$ are chosen from \cite{Claret:2004p807} to match the $u$ -band limb-darkening of the host star. |
We assume that the material absorbs or scatters radiation out of the line-of-sight with no scattering into the line-of-sight. which is valid for the optical depths required to produce the early-ingress transits. | We assume that the material absorbs or scatters radiation out of the line-of-sight with no scattering into the line-of-sight, which is valid for the optical depths required to produce the early-ingress transits. |
For this letter we assume the bow shock is of uniform density and that the material is static. however. our models are very general and can incorporate any density structure: analytic. tabulated. or from dynamical simulations. | For this letter we assume the bow shock is of uniform density and that the material is static, however, our models are very general and can incorporate any density structure: analytic, tabulated, or from dynamical simulations. |
Our goal is to determine the range of shock geometries that can provide both an early-ingress and suthcient optical depth in MgIT | Our goal is to determine the range of shock geometries that can provide both an early-ingress and sufficient optical depth in MgII |
absorption features present in the wavelength range centered on the Meg line triplet. (AA5164.5173.X... see. Tab. 7)) | absorption features present in the wavelength range centered on the Mg line triplet $\lambda\lambda\,5164,5173,5184$, see Tab. \ref{tab:TW_values}) ) |
using the Fourier Correlation Quotient 5184.method (Bender 1990: Bender et al. | using the Fourier Correlation Quotient method (Bender 1990; Bender et al. |
1994). as done in Paper I. We adopted LR 6817 (NILE) as the kinematical template to measure the stellar kinematics of NGC 3412. LR 7429 (IX31HE). for the kinematics of ESO 139-C009 and LC S74. and LR 3145 (ΝΕΤ) for the kinematics of NGC 1308 and NGC 1440. | 1994), as done in Paper I. We adopted HR 6817 (K1III) as the kinematical template to measure the stellar kinematics of NGC 3412, HR 7429 (K3III) for the kinematics of ESO 139-G009 and IC 874, and HR 3145 (K2III) for the kinematics of NGC 1308 and NGC 1440. |
The values of line-ol-sieht velocity. e... and. velocity dispersion c measured along the different slits for cach sample galaxyare given in Table 6.. | The values of line-of-sight velocity $v$, and velocity dispersion $\sigma$ measured along the different slits for each sample galaxyare given in Table \ref{tab:kinematics}. . |
penetrate so far. | penetrate so far. |
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